Solar Orbiter Science Activity Plan

ESA UNCLASSIFIED – For Official Use

Prepared by Yannis Zouganelis, Anik De Groof, Andrew Walsh, David Williams, Daniel Müller

Reference SOL-EST-PL-8539

Issue 0

Revision 1

Date of Issue 10 July 2017

Status Draft

Document Type

Distribution

Solar Orbiter Science Activity Plan


Page 2/312 Date 10 Jul 2017- Issue 0 Rev 1

Title

Issue 0 Revision

Author Yannis Zouganelis Date

Approved by Date

Daniel Müller (ESA Project Scientist)

Chris St. Cyr (NASA Project Scientist)

Javier Rodríguez-Pacheco (EPD PI)

Pierre Rochus (EUI PI)

Timothy Horbury (MAG PI)

Ester Antonucci (METIS PI)

Sami Solanki (PHI PI)

Milan Maksimovic (RPW PI)

Russ Howard (SoloHI PI)

Frédéric Auchère (SPICE Operations PI)

Sam Krucker (STIX PI)

Christopher J. Owen (SWA PI)

Reason for change Issue Revision Date

Issue 0 Revision

Reason for change Date Pages Paragraph(s)



Table of Contents:

1 INTRODUCTION ....................................................................................................................... 7 1.1 Purpose and Scope ........................................................................................................................ 7 1.2 Structure of the Document ........................................................................................................... 7 1.3 Assumptions and Definitions ....................................................................................................... 7 1.4 Applicable Documents ................................................................................................................. 8 1.5 Reference Documents ................................................................................................................... 8 2 MISSION PLANNING OVERVIEW ........................................................................................ 9 2.1 Mission Planning Cycle ................................................................................................................ 9 2.1.1 Long-Term Planning (LTP) ...................................................................................................... 10 2.1.2 Medium-Term Planning (MTP)................................................................................................ 10 2.1.3 Short-Term Planning (STP) ...................................................................................................... 10 2.1.4 Very-Short-Term planning (VSTP) .......................................................................................... 10 3 SCIENCE OBJECTIVES ......................................................................................................... 12 3.1 Solar Orbiter’s Top-Level Science Objectives ........................................................................... 12 3.1.1 What drives the solar wind and where does the coronal magnetic field originate? ................. 12 3.1.2 How do solar transients drive heliospheric variability? ........................................................... 13 3.1.3 How do solar eruptions produce energetic particle radiation that fills the heliosphere? .......... 13 3.1.4 How does the solar dynamo work and drive connections between the Sun and the heliosphere?

14 3.2 Lower-Level Objectives and Required Observations................................................................. 15 3.2.1 Objective 1: What drives the solar wind and where does the coronal magnetic field originate?

15 3.2.1.1 1.1 What are the source regions of the solar wind and heliospheric magnetic field? ............ 16 3.2.1.2 1.2 What mechanisms heat the corona and heat and accelerate the solar wind? ................... 32 3.2.1.3 1.3 What are the sources of solar wind turbulence and how does it evolve? ......................... 39 3.2.2 Objective 2: How do solar transients drive heliospheric variability? ....................................... 45 3.2.2.1 2.1 How do CMEs evolve through the corona and inner heliosphere? .................................. 45 3.2.2.2 2.2 How do CMEs contribute to solar magnetic flux and helicity balance? .......................... 49 3.2.2.3 2.3 How and where do shocks form in the corona? ............................................................... 51 3.2.3 Objective 3: How do solar eruptions produce energetic particle radiation that fills the

heliosphere? ....................................................................................................................................... 55 3.2.3.1 3.1 How and where are energetic particles accelerated at the Sun? ....................................... 55 3.2.3.2 3.2 How are energetic particles released from their sources and distributed in space and

time? 68 3.2.3.3 3.3 What are the seed populations for energetic particles? .................................................... 73 3.2.4 Objective 4: How does the solar dynamo work and drive connections between the Sun and the

heliosphere? ....................................................................................................................................... 76 3.2.4.1 4.0 Overall remarks and feasibility concerning Objective 4 observations with Solar Orbiter

77 3.2.4.2 4.1 How is magnetic flux transported to and re-processed at high solar latitudes? ............... 80 3.2.4.3 4.2 What are the properties of the magnetic field at high solar latitudes? ............................. 83 3.2.4.4 4.3 What is the nature of magnetoconvection? ...................................................................... 85 3.2.4.5 4.4 Are there separate dynamo processes acting in the Sun? ................................................. 86



3.2.4.6 4.5 How are coronal and heliospheric phenomena related to the solar dynamo? .................. 87 3.2.5 5. Additional science objectives ............................................................................................... 88 4 INSTRUMENT DESCRIPTIONS AND THEIR OPERATIONAL CONSTRAINTS ....... 93 4.1 EPD ............................................................................................................................................ 93 4.1.1 EPD Observables ..................................................................................................................... 93 4.1.1.1 Energy range .......................................................................................................................... 94 4.1.1.2 Mass resolution ...................................................................................................................... 95 4.1.1.3 Fields of view and angular resolution .................................................................................... 95 4.1.1.4 Geometric factor..................................................................................................................... 96 4.1.2 EPD modes ............................................................................................................................... 97 4.1.2.1 EPD Normal mode ................................................................................................................. 97 4.1.2.2 EPD Burst mode ..................................................................................................................... 97 4.2 EUI ............................................................................................................................................. 98 4.2.1 EUI observables ........................................................................................................................ 98 4.2.2 EUI modes and telemetry ......................................................................................................... 98 4.2.2.1 FSI Beacon mode (B) ........................................................................................................... 100 4.2.2.2 FSI Synoptic mode (S) ......................................................................................................... 100 4.2.2.3 FSI Reference Synoptic mode (R) ....................................................................................... 101 4.2.2.4 FSI Global eruptive event mode (G) .................................................................................... 102 4.2.2.5 FSI Find Event mode (FE) ................................................................................................... 103 4.2.2.6 FSI Faint High Corona mode (FHC) .................................................................................... 103 4.2.2.7 EUV & LYA Beacon modes (HB)....................................................................................... 104 4.2.2.8 EUV & LYA Coronal hole modes (C) ................................................................................. 104 4.2.2.9 EUV & LYA Quiet Sun modes (Q) ..................................................................................... 105 4.2.2.10 EUV & LYA Active Region modes (A) ......................................................................... 106 4.2.2.11 EUV & LYA Eruptive Event modes (E) ........................................................................ 106 4.2.2.12 EUV & LYA Discovery modes (D) ................................................................................ 107 4.3 MAG ......................................................................................................................................... 108 4.3.1 MAG observables ................................................................................................................... 109 4.3.2 MAG modes ........................................................................................................................... 109 4.3.2.1 MAG Normal mode ............................................................................................................. 109 4.3.2.2 MAG burst mode.................................................................................................................. 109 4.4 METIS ...................................................................................................................................... 110 4.4.1 METIS observables ................................................................................................................ 110 4.4.2 METIS modes ......................................................................................................................... 111 4.5 PHI ............................................................................................................................................ 119 4.5.1 PHI observables ...................................................................................................................... 119 4.5.2 PHI modes .............................................................................................................................. 119 4.6 RPW ......................................................................................................................................... 127 4.6.1 RPW observables .................................................................................................................... 129 4.6.2 RPW modes ............................................................................................................................ 129 4.7 SoloHI ...................................................................................................................................... 131 4.7.1 SoloHI observables ................................................................................................................. 131 4.7.2 SoloHI modes ......................................................................................................................... 131 4.8 SPICE ....................................................................................................................................... 134



4.8.1 SPICE observables ................................................................................................................. 134 4.8.2 SPICE modes .......................................................................................................................... 134 4.9 STIX ......................................................................................................................................... 145 4.9.1 STIX observables ................................................................................................................... 146 4.9.2 STIX modes ............................................................................................................................ 146 4.10 SWA ......................................................................................................................................... 148 4.10.1 SWA observables ............................................................................................................... 149 4.10.2 SWA modes ....................................................................................................................... 149 5 SCIENCE ACTIVITIES ......................................................................................................... 153 5.1 Introduction and SOOPs ........................................................................................................... 153 5.2 List of SOOPs ........................................................................................................................... 153 5.2.1 I_DEFAULT ........................................................................................................................... 154 5.2.2 L_IS_STIX ............................................................................................................................... 161 5.2.3 L_IS_SoloHI_STIX ................................................................................................................ 163 5.2.4 L_FULL_LRES_MCAD_Coronal_Synoptic ......................................................................... 166 5.2.5 L_FULL_LRES_MCAD_ProbeQuadrature ........................................................................... 169 5.2.6 L_FULL_MRES_MCAD_CME_SEPs .................................................................................. 171 5.2.7 L_FULL_HRES_LCAD_MagnFieldConfig .......................................................................... 174 5.2.8 L_FULL_HRES_MCAD_Coronal_He_Abundance .............................................................. 176 5.2.9 L_FULL_HRES_HCAD_Eruption_Watch ............................................................................ 179 5.2.10 L_FULL_HRES_HCAD_Coronal_Dynamics................................................................... 182 5.2.11 L_SMALL_MRES_MCAD_Ballistic-connection............................................................. 186 5.2.12 L_SMALL_MRES_MCAD_Connection_Mosaic............................................................. 188 5.2.13 L_SMALL_HRES_HCAD_Fast_Wind............................................................................. 192 5.2.14 L_SMALL_HRES_HCAD_SlowWindConnection........................................................... 197 5.2.15 L_BOTH_LRES_MCAD_Pole-to-Pole............................................................................. 204 5.2.16 L_BOTH_MRES_MCAD_Farside_Connection ............................................................... 206 5.2.17 L_BOTH_MRES_MCAD_Flare_SEPs ............................................................................. 209 5.2.18 L_BOTH_HRES_LCAD_CH_Boundary_Expansion ....................................................... 215 5.2.19 R_FULL_LRES_HCAD_GlobalHelioseismology ............................................................ 219 6 PLANNING STRATEGY ....................................................................................................... 252 7 OCTOBER 2018 OPTION E TRAJECTORY AND MEDIUM TERM PLANNING...... 267 7.1 Option E trajectory ................................................................................................................... 267 7.2 Planning periods for Option E (MTPs) .................................................................................... 267 7.2.1 MTP05 - 2021/01/01 - 2021/07/01 ......................................................................................... 268 7.2.1.1 RS window (default) placement ........................................................................................... 268 7.2.1.2 Science planning .................................................................................................................. 268 7.2.2 MTP06 - 2021/07/01 - 2022/01/01 ......................................................................................... 270 7.2.2.1 RS windows (default placement): ........................................................................................ 270 7.2.2.2 New placement for the RS window: .................................................................................... 270 7.2.2.3 SOOP planning..................................................................................................................... 271 7.2.3 MTP07 - 2022/01/01 - 2022/07/01 ......................................................................................... 273 7.2.3.1 RS window (default) placement ........................................................................................... 273 7.2.3.2 Current scenario ................................................................................................................... 273 7.2.4 MTP08 - 2022/07/01 - 2023/01/01 ......................................................................................... 276



7.2.4.1 RS window (default) placement ........................................................................................... 276 7.2.4.2 Current scenario ................................................................................................................... 276 7.2.5 MTP09 - 2023/01/01 - 2023/07/01 ......................................................................................... 279 7.2.6 MTP10 - 2023/07/01 - 2024/01/01 ......................................................................................... 281 7.2.7 MTP11 - 2024/01/01 - 2024/07/01 ......................................................................................... 284 7.2.8 MTP12 - 2024/07/01 - 2025/01/01 ......................................................................................... 287 7.2.9 MTP13 - 2025/01/01 - 2025/07/01 (EMP) ............................................................................. 289 7.2.10 MTP14 - 2025/07/01 - 2026/01/01 .................................................................................... 291 7.2.11 MTP15 - 2026/01/01 - 2026/07/01 .................................................................................... 293 7.2.12 MTP16 - 2026/07/01 - 2027/01/01 .................................................................................... 295 7.2.13 MTP17 - 2027/01/01 - 2027/07/01 .................................................................................... 297 7.2.14 MTP18 - 2027/07/01 - 2028/01/01 .................................................................................... 299 7.2.15 MTP19 - 2028/01/01 - 2028/07/01 .................................................................................... 301 7.2.16 MTP20 - 2028/07/01 - 2029/01/01 .................................................................................... 303 7.2.17 MTP21 - 2029/01/01 - 2029/07/01 .................................................................................... 304 8 SIMULATIONS ....................................................................................................................... 306 9 APPENDIX .............................................................................................................................. 307 9.1 Bibliography ............................................................................................................................. 307


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1 INTRODUCTION

1.1 Purpose and Scope

The Science Activity Plan (SAP) for Solar Orbiter describes in a structured way all scientific

activities to be carried out by the instruments throughout all science mission phases in order to fulfil

the Science Requirements of the mission. It tracks how high-level science objectives are mapped to

more specific scientific objectives and these, in turn, to scientific activities that will be scheduled at

specific times during the mission. As context information, it also contains instrument operations

scenarios and modelling, and a description of all mission phases.

This first draft, which is not yet approved by the SWT, is issued as Reference Document for the

Solar Orbiter OGS-SGS Joint Ground Segment Implementation Review. It has been built by the

joint efforts of the Science Working Team (SWT), the Science Operations Working Group

(SOWG), instrument teams, PIs and Co-Is, and external experts, under the coordination and

guidance of the Science Operations Centre (SOC) team and the Project Scientists.

In general, it is foreseen to make use of the successful Joint Operations Plan (JOP) concept of the

SOHO mission, which is also being used by the Hinode mission, adapted to the specific needs of

Solar Orbiter (hence to be called SOOP). The general characteristics of the scientific mission

planning approach for Solar Orbiter are described in Section 2.

1.2 Structure of the Document

Section 2 provides an overview of the mission planning process. Section 3 describes the mission’s

science objectives, on the basis of [AD01, AD02]. Section 4 contains brief descriptions of each

instrument along with its operational modes and constraints. Science activities, defined as sets of

observations (SOOPs) which address each of the scientific objectives, are described in Section 5.

Section 6 describes the planning strategy that is used for filling the mission timeline with the various

science activities. Section 7 provides mission profiles (orbits etc.) and the description of all

Medium-Term Planning periods together with the activities planned in each one. Section 8 will

contain in the future the results of the simulations performed regarding the SSMM fill state, the

instrument stores fill state, power consumption etc. Any additional information, as well as a

bibliography of scientific papers is provided in Section 9.

1.3 Assumptions and Definitions

This first draft (SAP v0) is sufficiently detailed for its purpose as a reference document for the

Ground Segment Implementation Review. However, it still lacks information and details before the

first version can be released. The following assumptions and limitations will be relaxed in future

versions:

• The Cruise Phase is not yet described.

• An October 2018 launch date has been assumed with the Option E trajectory. This is only

affecting the Sections 7-8 and part of Section 6. Sections 3-5 are trajectory independent.

• Even though the Objective 4 sub-objectives have been well detailed and the corresponding

SOOPs have been described, these have not been planned in the mission timeline. This is



due to the fact of the recent project delays and postponement of the launch date, which

resulted in equivalent delays of the SAP activities. These will be resumed once the new

launch date and information on trajectory get stabilized.

• For the same reasons, the preliminary simulations (Section 8) performed for the current

version of the SAP are not included.

1.4 Applicable Documents

[AD01] Solar Orbiter Science Management Plan (SMP, SOL-EST-PL-00880)

[AD02] Solar Orbiter Science Requirements Document (SciRD, SOL-EST-RS-1858)

1.5 Reference Documents

[RD01] Technical Note: Solar Orbiter Pointing Strategy during Remote-Sensing Windows: Science

Requirements (SOL-EST-TN-4020)

[RD02] Solar Orbiter – Exploring the Sun-Heliosphere Connection (Solar Physics, Vol. 285, Issue

1-2, pp. 25-70, 2013)

[RD03] Solar Orbiter Science Implementation Requirements Document (SIRD, SOL-EST-RS-

4514)

[RD04] Solar Orbiter Science Operations Centre Science Implementation Plan (SIP, SO-SGS-PL-

001)

[RD05] Solar Orbiter Science Operations Concept Document (SOL-SGS-PL-0001)

[RD06] Solar Orbiter Instrument Operation Request Interface Control Document (IOR ICD, SOL-

SGS-ICD-0003)



2 MISSION PLANNING OVERVIEW

While the general mission planning approach for all routine science operations of Solar Orbiter will

be built on the experience of ESA’s precursor solar system missions Mars Express, Venus Express,

and Rosetta, a fundamental difference with respect to planetary missions is the highly dynamic

nature of the Sun. Given the short time-scales on which the targets of remote-sensing observations

(e.g. solar active regions) change, together with the narrow fields-of-view of the high-resolution

imaging telescopes (which cover less than 3% of the solar disk at perihelion), turn-around times

between defining the pointing and executing the observations of at most three days is required

during the Remote-Sensing Windows, as described in [RD01].

The resulting requirements on the science operations planning for Remote-Sensing Windows are

summarized below, on the basis of [RD03]. The mission planning cycle described is the result of a

collaborative effort of OGS, SOC and PS.

2.1 Mission Planning Cycle

For each mission phase (Cruise Phase (CP), Nominal Mission Phase (NMP), Extended Mission

Phase (EMP)), a baseline science plan will be established and documented in the SAP before each

mission phase commences. This plan will take into account the general characteristics and major

constraints of each orbit. The SAP will be a ‘live’ document in the sense that it will be frequently

updated as more and more science activities get worked out in detail, feedback from earlier

observations (and their planning) is injected into the planning of later orbits, etc.

For each orbit, a typical mission planning cycle starts with the SWT deciding upon the top-level

science objectives for this orbit, based on the general goals previously agreed upon by the SWT and

defined in the SAP. Given this input, the SOWG defines a coherent mission-level observing plan. In

this task, they will be assisted by the SOC, which will provide detailed information on the resources

available (e.g., on-board memory management, telemetry downlink). In turn, the PI teams provide,

at fixed deadlines and with a fixed periodicity, inputs to the SOC for the requested science

operations, which implement the observing plan defined by the SOWG. The SOC passes a

consolidated request to the MOC which checks the requests against mission, environmental and

resource constraints.

It is important to stress that top-level science operations planning for the mission needs to be done

well in advance due to the severe constraints on Solar Orbiter’s data downlink volume and on-board

storage in the Solid State Mass Memory (SSMM). In particular, this entails feasibility studies of

planned operations taking into account the expected time-dependent telemetry downlink profile as

well as SSMM load levels.

The mission planning cycle for the routine science operations phase will therefore be divided into

different levels, as described in RD05:



2.1.1 Long-Term Planning (LTP)

LTP will fix the ground station schedule. This process typically will take place once for each

mission phase (covering the entire mission phase) until the start of the NMP, and once for every 6-

months period, approximately covering one orbit around the Sun, after the start of NMP. In each

orbit, this process will in general commence after the end of the last Remote-Sensing Window,

counted from aphelion. At LTP level, event skeleton fields are provided by the OGS to the SGS,

which will contain all the information required by the SGS to conduct science planning.

2.1.2 Medium-Term Planning (MTP)

MTP will fix the usage of spacecraft resources. This process will typically take place at several

intervals during Cruise Phase, and once per orbit during NMP, covering the entire orbit. At MTP

level, the top-level science operations plan for the entire orbit is defined and a default pointing

profile for Remote-Sensing Windows is defined (with possible updates at VSTP level, see below).

2.1.3 Short-Term Planning (STP)

STP will generate detailed schedules of commands for the spacecraft and for the ground stations.

This process will take place typically every week covering one week as described in detail in RD06.

At STP level, the instrument activities can be modified, provided they fit into the resource envelope

defined at MTP level.

2.1.4 Very-Short-Term planning (VSTP)

In the case of Remote Sensing Windows (RSWs), VSTP may be required. This planning level, with

turn-around times of at most three days between observations and execution of the new Pointing

Request (PTR), is required for RSWs in which features on the solar disk, e.g. active regions, shall be

tracked over time. This is due to the short lifetimes and non-deterministic motion of targets on the

Sun [RD01] and the relatively large (compared to the fields-of-view of the high-resolution imaging

telescopes) and temperature-dependent absolute pointing error (APE) of the spacecraft. This VSTP

consists of (i) initial target selection and (ii) updates to the pointing.

Prior to the start of an RSW, a limited set of precursor observations with the full-disk imaging

telescopes of the EUI and PHI instruments is performed and downlinked with high priority. Based

on the returned data, the target for the start of the RSW will be defined. This step is required to

make a decision on the pointing of the spacecraft and, in turn, the high-resolution imaging

telescopes. In case the orbital constellation permits making this decision by means of other

observations (e.g. using ground-based telescopes), this step can be omitted.

During the course of an RSW, a limited set of daily low-latency data, consisting of full-disk and

high-resolution images, will be downlinked with high priority. Based on the evaluation of these

images, the pointing will be updated by means of uploading a PTR. The final science target

selection (and retargeting) shall be responsibility of a person nominated by the SOWG, who shall

make him/herself available according to the schedule required by the planning process. Refinement



of the tracking of this target shall be done by SOC, according to a set of rules established by the

SOWG.



3 SCIENCE OBJECTIVES

Solar Orbiter’s mission is to address the central question of heliophysics: How does the Sun create

and control the heliosphere? This, in turn, is a fundamental part of the second science question of

ESA’s Cosmic Vision programme: “How does the solar system work?” Solar Orbiter is specifically

designed to identify the origins and causes of the solar wind, the heliospheric magnetic field, solar

energetic particles, transient interplanetary disturbances, and the Sun's magnetic field itself.

The supersonic solar wind, driven by dynamic plasma and magnetic processes at the Sun’s surface,

expands to surround the solar system’s planets and the space far beyond. Below the surface, the

solar dynamo drives magnetic fields whose buoyancy brings them to the surface where they form

huge arcades of loops, which contain enormous amounts of stored energy. These magnetic loops are

stretched and sheared by the Sun’s differential rotation and unknown surface processes, eventually

erupting in explosions, which eject magnetic structures that fly into the solar system, occasionally

impacting the Earth and its magnetic shield with disruptive effects on space and terrestrial systems.

Understanding the complex physical processes at work in this system is the central goal of

heliophysics. Since the Sun and presumably the heliosphere are typical of many small stars and their

stellar spheres, these studies are relevant to astrophysics, but are unique since the Sun alone is close

enough for detailed study.

Over the past ~20 years, an international effort to understand the Sun and heliosphere has been

undertaken with an array of spacecraft carrying out both remote observations at visible, UV, and X-

ray wavelengths, as well as in-situ observations of interplanetary plasmas, particles, and fields.

Combined and coordinated observations from missions such as Ulysses, Yohkoh, SOHO, TRACE,

RHESSI, Hinode and STEREO have resulted in an enormous advance in our understanding of the

Sun and heliosphere, and have proven that critical progress in understanding the physics requires

both remote and in-situ observations working together.

Although our vantage point at 1 AU is close by astrophysical measures, it has been long known that

much of the crucial physics in the formation and activity of the heliosphere takes place much closer

to the Sun, and that by the time magnetic structures, shocks, energetic particles and solar wind pass

by Earth they have already evolved and in many cases mixed so as to blur the signatures of their

origin. With the proven effectiveness of combined remote and in-situ studies on the missions cited

above, it is expected that critical new advances will be achieved by flying a spacecraft combining

remote and in-situ observations into the inner solar system. From this inner-heliospheric vantage

point, solar sources can be identified and studied accurately and combined with in-situ observations

of solar wind, shocks, energetic particles, etc., before they evolve significantly.

3.1 Solar Orbiter’s Top-Level Science Objectives

The four top-level scientific questions to be addressed by Solar Orbiter are:

3.1.1 What drives the solar wind and where does the coronal magnetic field originate?

The solar corona continuously expands and develops into a supersonic wind that extends outward,

interacting with itself and with the Earth and other planets, to the heliopause boundary with

interstellar space, far beyond Pluto’s orbit. The solar wind has profound effects on planetary

environments and on the planets, themselves – for example, it is responsible for many of the



phenomena in Earth’s magnetosphere and is thought to have played a role in the evolution of Venus

and Mars through the erosion of their upper atmospheres.

Two classes of solar wind – ‘fast’ and ‘slow’ – fill the heliosphere, and the balance between them is

modulated by the 11-year solar cycle. The fast solar wind (~700km/s and comparatively steady) is

known to arise from coronal holes. The slow solar wind (~400 km/s) permeates the plane of the

ecliptic during most of the solar cycle so it is important to Earth's space environment. The slow solar

wind differs from the fast wind in mass flux and composition, which is consistent with confined

plasma in the solar corona. The specific escape mechanism through the largely closed magnetic field

is not known since candidate sites and mechanisms cannot be resolved from 1 AU. Fast and slow

wind carry embedded turbulent fluctuations, and these also display different properties compatible

with different solar origins. It is thought that such fluctuations may be responsible for the difference

in heating and acceleration between different solar wind streams.

Understanding the physics relating the plasma at the solar surface and the heating and acceleration

of the escaping solar wind is crucial to understanding both the effects of the Sun on the heliosphere

and how stars in general lose mass and angular momentum to stellar winds.

3.1.2 How do solar transients drive heliospheric variability?

The largest transient events from the Sun are coronal mass ejections (CMEs), large structures of

magnetic field and material that are ejected from the Sun at speeds up to 3000 km/s. CMEs are also

of astrophysical interest since they are the dominant way that stars shed both magnetic flux and

magnetic helicity that build up as a result of the stellar dynamo. Interplanetary CMEs (ICMEs) are

the major cause of interplanetary shocks, but the locations and mechanisms by which shocks form

around them are not known since they occur in the inner solar system. Similarly, the longitudinal

structure of ICMEs is not directly observable from the ecliptic, while their extent has a large impact

on the acceleration of energetic particles. ICMEs are a major cause of geomagnetic storms but their

effectiveness at disrupting the magnetosphere is only loosely related to the parent CME, because the

evolution of the propagating cloud with the surrounding heliosphere is complex and has not been

well studied. These unknowns have direct impact on our ability to predict transient (“space

weather”) events that affect Earth.

3.1.3 How do solar eruptions produce energetic particle radiation that fills the

heliosphere?

Like many astrophysical systems, the Sun is an effective particle accelerator. Large solar energetic

particle (SEP) events produce highly energetic particles that fill the solar system with ionizing

radiation. CME-driven shocks can produce relativistic particles on time scales of minutes, and many

CMEs convert ~10% of their kinetic energy into energetic particles. Other processes produce high-

energy particles on magnetic loops without involving shocks. The multiple processes operating in

SEP events are not well understood or distinguishable from observations at 1 AU. In particular,

particles accelerated in the corona and inner heliosphere are scattered by inhomogeneities in the

interplanetary magnetic field (IMF) before they arrive at Earth, destroying much of the information

they carry about the processes that accelerated them. Particle transport and scattering in the inner

solar system are poorly understood since the turbulence properties cannot be determined from 1 AU.

The actual seed population of particles energized by CME-driven shocks in the inner solar system is



unexplored, and needs to be understood to construct a complete picture of particle acceleration in

shock-related events.

3.1.4 How does the solar dynamo work and drive connections between the Sun and the

heliosphere?

The Sun’s magnetic field connects the interior of the star to interplanetary space and is dominated

by a quasi-periodic 11-year sunspot cycle that modulates the form of the heliosphere and strongly

affects the space environment throughout the solar system. The large-scale solar field is generated in

the Sun’s interior, within the convection zone, by a dynamo driven by complex three-dimensional

mass flows that transport and process magnetic flux. Despite notable advances in our knowledge

and understanding of solar magnetism made possible by Ulysses, SOHO, and Hinode observations

as well as by recent theoretical models and numerical simulations, fundamental questions remain

about the operation of the solar dynamo and the cyclic nature of solar magnetic activity. Of

paramount importance to answering these questions is detailed knowledge of the transport of flux at

high latitudes and the properties of the polar magnetic field. To date, however, the solar high

latitudes remain poorly known owing to our dependence on observations made from the ecliptic. In

addition to questions about the global dynamo and the generation of the large-scale field, there are

unanswered questions about the origin of the small-scale internetwork field observed in the quiet

photosphere. Is this weak field produced by turbulent local dynamo action near the solar surface?



3.2 Lower-Level Objectives and Required Observations

In this section, we describe detailed lower-level objectives of each of the above four top science

objectives.

3.2.1 Objective 1: What drives the solar wind and where does the coronal magnetic field

originate?

Hot plasma in the Sun’s atmosphere flows radially outward into interplanetary space to form the

solar wind, filling the solar system and blowing a cavity in the interstellar medium known as

the heliosphere. During solar minimum, large-scale regions of a single magnetic polarity in the

Sun’s atmosphere – polar coronal holes – open into space and are the source of high speed (~700

km/s), rather steady solar wind flows. There is also a slow wind (300-500 km/s) that emanates from

magnetically complex regions at low latitudes and the periphery of coronal holes. It is highly

variable in speed, composition, and charge state. The origin of the slow wind is not known. At solar

maximum, this stable bimodal configuration gives way to a more complex mixture of slow and fast

streams emitted at all latitudes, depending on the distribution of open and closed magnetic regions

and the highly tilted magnetic polarity inversion line.

The fast wind from the polar coronal holes carries magnetic fields of opposite polarity into the

heliosphere, which are then separated by the heliospheric current sheet (HCS) embedded in the slow

wind. Measurements over a range of latitudes far from the Sun show that this boundary is not

symmetric around the Sun’s equator, but is on average displaced southward. This offset must reflect

an asymmetry on the Sun; but since there cannot be a mismatch between the inward and outward

magnetic flux on the Sun, its origin is unclear. In situ, the HCS is warped and deformed by the

combined effects of solar rotation and inclination of the Sun’s magnetic axis, effects that are even

more prominent at solar maximum.

The energy that heats the corona and drives the wind comes from the mechanical energy of

convective photospheric motions, which is converted into magnetic and/or wave energy. In

particular, both turbulence and magnetic reconnection are implicated theoretically and

observationally in coronal heating and acceleration. However, existing observations cannot

adequately constrain these theories, and the identity of the mechanisms that heat the corona and

accelerate the solar wind remains one of the unsolved mysteries of solar and heliospheric physics.

How the coronal plasma is generated, energized, and the way in which it breaks loose from the

confining coronal magnetic field are fundamental physical questions with crucial implications

for predicting our own space environment, as well as for the understanding of the natural plasma

physics of other astrophysical objects, from other stars, to accretion disks and their coronae, to

energetic phenomena such as jets, X- and gamma-ray bursts, and cosmic-ray acceleration.

The solar wind contains waves and turbulence on scales from millions of kilometers to below the

electron gyroradius. The turbulence scatters energetic particles, affecting the flux of particles that

arrives at the Earth; local kinetic processes dissipate the turbulent fluctuations and heat the plasma.

Properties of the turbulence vary with solar wind stream structure, reflecting its origins near the

Sun, but the turbulence also evolves as it is carried into space with the solar wind, blurring the

imprint of coronal conditions and making it difficult to determine its physical origin. The inner

heliosphere, where Solar Orbiter will conduct its combination of remote-sensing and in-situ



observations, provides the ideal laboratory for understanding the magnetohydrodynamic turbulence

of natural plasmas expected to be ubiquitous in astrophysical environments.

In the following sections, we discuss in more detail three interrelated questions which flow down

from this top-level question: What are the source regions of the solar wind and the heliospheric

magnetic field? What mechanisms heat and accelerate the solar wind? What are the sources of

turbulence in the solar wind and how does it evolve?

3.2.1.1 1.1 What are the source regions of the solar wind and heliospheric magnetic field?

1.1.1 Source regions of the fast solar wind

1.1.1.1 Low FIP fast wind origins

Description of the objective:

The fast wind, which does not exhibit strong first ionization potential (FIP) enhancements, could

come directly from the photosphere, from small cool coronal loops and open magnetic funnels (Tu

et al., 2005; Schwadron and McComas, 2003) at the base of coronal holes or spicules, which also

exhibit small FIP enhancements. Remote observations have revealed many cases of macrospicules

undergoing reconnection and erupting within coronal holes (Yamauchi et al., 2005). Do they

contribute to the fast solar wind streams? Polar plumes have also long been suspected to be a

significant source of fast solar wind (Deforest et al., 1997), as well as polar regions within plumes

("interplume lanes") (Giordano et al., 2000). Micro-streams of plasma originating in the coronal

holes may be related to polar plumes (Neugebauer et al., 1995), though evidence for this is

controversial (e.g., McComas et al., 1996). However, the relation could be difficult to observe with

Solar Orbiter since large amplitude Alfvénic fluctuations generate micro-streams signals in the fast

stream (Matteini et al., 2013).

Remarks:

In order to address this objective, we need to observe if the fast wind comes from the different

above sources. This objective can be split into two different observational strategies:

• A pure coronal characterization by pointing to the center of a coronal hole and its

boundaries, mostly focused on remote sensing observations (closest possible at high

latitude).

• An attempt to observe fast wind in situ and link it to its origin back to the coronal hole. In

order to be connected, we have to choose a coronal hole at the west limb, and observe it

from a location close to the Sun (at least one perihelion window would be good for the high-

resolution observations of plume-interplume FIP variation) and preferably during high-

latitude windows. Another possibility would be to observe during a window before



perihelion in order to obtain the overall context and then observe during the perihelion for

connectivity. Polar coronal holes could also be observed with METIS from the equatorial

plane (no need of high-latitude windows, since it can observe the sky above the hole). In this

case, however, there is less chance of doing linkage science. Low-latitude coronal holes

should also be observed for intermediate speed outflow. Even if this has a lower priority, it

should not be neglected since this is the typical solar wind observed at Earth.

• The minimum of the solar activity cycle is preferable and/or the declining phase for the polar

coronal holes. Low-latitude coronal holes can be observed at any phase of the cycle, but the

probability to be at the right longitude at perihelion is low.

The needed observations include:

• EUI/HRI in Coronal Hole mode with 1 min of cadence, during 1-2 hours and 12 hours at

lower cadence.

• PHI at high resolution at 1 min cadence.

• SPICE: FIP and velocity maps. A combination of composition, dynamics and a 30’’-wide

movie. 3.2 hours for SPICE, 3 times per day.

• In situ instruments: normal mode for the connection and burst modes for more details. MAG

Burst mode is required for ion cyclotron wave identification in order to distinguish from

different acceleration and heating mechanisms as well as small scale changes. Recent Helios

data analysis shows velocity changes with jet-like features probably well below the 40s

cadence. We would, therefore, need 3D distributions at 1s scales in order to determine the

properties inside and outside these features.

Other remarks:

• Additional observations from near-Earth assets are desirable but not required.

• EMC Quiet is required for linkage science. Noisy periods can be tolerated during the RS

observations, but EMC quiet is required approximately 12 hours later (depending on which

distance Solar Orbiter is).

• EPD, STIX are not required for this objective.

The possible remote sensing targets should include wide regions of well extended coronal holes, as

well as smaller regions for focusing on the following:

• Small cool coronal loops.

• Open magnetic funnels at the base of coronal holes (Tu et al., 2005; Schwadron and

McComas, 2003).

• Spicules with short lifetimes, fast motions, and hot plasma components. Macrospicules

reconnection and eruption within coronal holes (Yamauchi et al., 2005).

• Polar plumes (Deforest et al., 1997).

• Polar regions within plumes (“interplume planes”) (Giordano et al. 2000).

• Diverging polar regions of the extended corona.

• Coronal hole boundaries.



The SOOP that addresses this objective is L_SMALL_HRES_HCAD_Fast_Wind, which also

addresses objective 1.1.1.2.

1.1.1.2 Origin of the small-scale X-ray and UV jets in polar coronal holes


One of the biggest discoveries of the Hinode satellite is the frequent occurrence of small-scale X-ray

and UV jets in polar coronal holes (Cirtain et al. 2007; Nistico et al., 2009; Krucker et al., 2011).

Such jets have widths between 2 × 103 and 2 × 104 km. Their origin is thought to be magnetic

reconnection of coronal field lines. At the reconnection site, Alfvén waves develop and produce

outflow velocities up to 800 kms-1, while the energy released by the reconnection heats the plasma

locally, generating mass motions with sonic speeds of ∼200 kms-1.

Given the high velocities and frequency of these events, it has been suggested that they contribute to

the fast solar wind. Their relations to the photospheric magnetic field, or the relevance of

photospheric processes for triggering them, have not been established yet. The high latitudes at

which they are observed hamper the accurate determination of their photospheric footpoints either

from the ground or NEO. Observations with remote-sensing instruments aboard Solar Orbiter will

allow us to understand the relation of these jets from a unique vantage point, providing comparable

high-angular resolution data simultaneously in the corona, chromosphere, and photosphere.

Remarks:

This objective could also be considered as part of the previous one (1.1.1.1). It only appears in a

separate section because it mostly requires remote sensing observations including STIX. From the

operational point of view, it can be jointly addressed with 1.1.1.1 with the

SOOP L_SMALL_HRES_HCAD_Fast_Wind. Observation of a sufficiently extended coronal hole

is required in order to catch multiple jets.

1.1.2 Source regions of the slow solar wind

1.1.2.1 Does slow solar wind originate from the over-expanded edges of coronal holes

(Antonucci et al., 2006)?


Study of the anti-correlation between the speed of the solar wind and the expansion rate of the

magnetic field. Wang and Sheeley (1990), Arge and Pizzo (2000), and others observe an anti-

correlation between solar wind speed and the so-called expansion factor of a flux tube near the Sun

calculated with a Potential-Field-Source-Surface (PFSS) model. Near the edges of coronal holes, the

flux-tube expansion factor is larger and generates a slower solar wind (Antonucci et al., 2006). The

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/L_SMALL_HRES_HCAD_Fast_Wind




physical connection between flux-tube expansion and solar wind speed is unclear. Larger flux-tube

expansion increases the downward electron and proton heat flux (e.g., Cranmer et al., 2007) and

creates an energy sink that slows the solar wind (Schwadron and McComas, 2003). However,

detailed solar models (e.g., Lie-Svendsen et al., 2002) have difficulty accounting for the solar wind

properties under these circumstances. If slow solar wind has a coronal hole origin, we expect a

smooth transition across the coronal hole boundary. Boundary layers outside the coronal hole

(Schwadron et al., 2005) need to be correlated with their remotely observed source

structures. Furthermore, an “anomalous” slow solar wind is observed to be originating from the

boundaries of coronal holes with the same level of Alfvénic fluctuations as in the nearby fast wind,

but of considerable smaller amplitude (Antonucci et al., 2005; D’Amicis et al., 2015).

Remarks:

• This objective can be addressed with two combined observing strategies:

o At first, several (2-3) days of disc center pointing for overall coronal hole

configuration during a perihelion window that would scan over many latitudes. PHI

and EUI/FSI will give us a ~3D view of the coronal hole edges and METIS can

observe close to the Sun, while SPICE will do composition mapping scanning.

o At a second phase, the coronal hole boundary will be mapped with SPICE mosaics

during 1 day (with METIS being off). The SPICE raster area should be optimized to

make sure that both open and closed field boundary is captured. The choice of the

lines has to be optimized depending on the type of target.

• An alternative strategy would be to observe with SoloHI and METIS during a high-latitude

window, before trying to do the linkage science at the following perihelion window.

• The connectivity being of main importance for the solar wind origin objectives, we will need

modeling in order to choose the right pointing.

• Observations during two successive windows (high-latitude – especially for high-latitude

coronal holes - and the following perihelion) are preferred.

• It is easier to perform these observations during the minimum or declining phase of solar

activity, when we have a strong magnetic dipole. Even though it will not be easy to choose

where to point during the maximum, it is worth planning it as well at least once during the

maximum.

• Earth context observations are important for better constraining the models and improving

pointing decisions.

• Parker Solar Probe joint observations at a radial alignment would be beneficial, especially if

PSP happens to be at its perihelion.

The SOOP that addresses this objective is L_BOTH_HRES_LCAD_CH_Boundary_Expansion,

which was specifically created for this objective and does not currently cover any other sub-

objectives. All other slow wind origin objectives are addressed by the similar SOOP

L_SMALL_HRES_HCAD_SlowWindConnection.

1.1.2.2 Does slow and intermediate solar wind originate from coronal loops outside of coronal

holes?

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/L_BOTH_HRES_LCAD_CH_Boundary_Expansion

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/L_SMALL_HRES_HCAD_SlowWindConnection




Solar observations show that the elemental composition, temperatures and thus the charge states of

large coronal loops outside of coronal holes are similar to the composition of the slow solar wind

(Feldman et al. 2005; Baker et al., 2013; Brooks et al., 2015). Loop source models require that the

foot-points of open field lines are interspersed with large loops outside coronal holes (Fisk et al.,

1999; Fisk and Schwadron, 2001). Interchange reconnection between open field lines and loops

(along topological boundaries: quasi-separatrix layers, including coronal nulls, e.g. van Driel-

Gesztelyi et al., 2012) releases the material stored on the loops and generates the slow wind.

Observations indicate the legs or tips of the helmet streamer, loops near active regions, and 80-300

Mm loops may all contribute to the slow wind. The challenge is to associate definitively the features

and structures in the slow wind with the morphology of the coronal complex. A loop origin for the

slow wind would show sharp interfaces and characteristic variations that can be correlated to the

remotely observed sources. Other interesting aspects to be explored include the existence of narrow

regions of open field lines between multiple closed field configurations within the streamer belts

(Noci et al., 1997; Wang et al., 2000; Noci and Gavryuseva, 2007), slow wind coming from small

coronal holes near or inside active regions (Wang, 2017, 10.3847/1538-4357/aa706e) vs from open

magnetic flux rooted in active regions (Liewer et al., 2004, 10.1007/s11207-004-1105-z), the

differences with the solar activity maximum when the slow wind appears to emanate rather from

small coronal holes and active regions (e.g. Neugebauer et al., 2002), the possibility of turbulent

reconnection (Rapazzo et al., 2012), the quasi-steady flow from streamer legs (Suess and Nerney,

2006), the transient contribution to the slow wind from parcels of plasma escaping from the

streamer core (Suess et al., 2009), explore the sharp interfaces in the slow wind indicative of loop

origin, observe the fans at the edge of active regions.

Remarks:

• This objective corresponds to different kinds of structures (helmet streamer, loops near

active regions, streamer core, belts, edge of active regions...). These will be different targets

observed by the same SOOP.

• As for all solar wind origin objectives, we need to compare the elemental composition,

temperatures and charge states of large coronal loops to the in situ slow wind ones (Feldman

et al., 2005; Baker et al., 2013; Brooks et al., 2015).

• Several days of observations of a large region surrounding an emerging active region are

required, it is probably interesting to observe during a full window for studying the

abundance of active regions.

• Observations near the perihelion are preferred for an increased resolution and a higher

probability for connectivity.

http://dx.doi.org/10.3847/1538-4357/aa706e

http://dx.doi.org/10.1007/s11207-004-1105-z



• For the high-resolution observations, EUI needs to support SPICE with at least one or a few

high-resolution frames for each SPICE raster.

• As for all solar wind connectivity objectives, an EMC Quiet period is needed some hours

after the remote sensing observations. Note, however, that this is not easy to effectively plan

in practice since this would decrease the duration of the remote sensing observation of the

active region.

• Parker Solar Probe joint observations at a radial alignment would be beneficial, especially if

PSP happens to be at its perihelion.

• Reconnection events at the boundaries of a coronal hole caused by the emergence of a new

active region occur on timescales of the order of 1 hour (Rappazzo et al. 2012). Edmondson

et al. (2010) also found timescales of the order of 1-2 hours for interchange reconnection

processes between an active region loop and an adjacent coronal hole. The typical network

supergranule reconfiguration time is of the order of 1-2 days. This timescale can also be

adopted as the reconfiguration time of the global magnetic field lines. The magnetic dipole

emerging rate in coronal holes is 1-2/day (Abramenko et al. 2006). Schrijver at al. (1998)

found that flux concentrations are enhanced and disappear with a characteristic timescale of

about 1.5 days. Fisk and Schwadron (2001) state that the characteristic time for a change in

the open flux, due to reconnection with loops at low latitudes, is about 36-38 hr. On the other

hand, Antiochos et al. (2011) describe the large-scale field evolution as approximated to a

sequence of topologically smooth quasi-steady states (quasi-steady models). Therefore, we

are facing with phenomena occurring on timescales ranging from about 1 hour, if we

consider a single active region, to about 36 hr, when we consider the global coronal

magnetic field. On this basis, a PHI cadence of 6 hr is enough for investigating interchange

reconnection phenomena at a global scale (larger than the active region FoV). A shorter

cadence (about 1 hr) might be required to investigate at high spatial resolution interchange

reconnection phenomena between an active region and an adjacent coronal hole.

This objective is addressed by the SOOP L_SMALL_HRES_HCAD_SlowWindConnection. The

similar SOOP L_SMALL_MRES_MCAD_Connection_Mosaic can be used if we want to cover a

wider area around the active region with lower resolution, in particular for the SPICE FoV.

1.1.2.3 Abundance of minor ions as a function of height in the corona as an indicator of slow

or fast wind.


Abundance of minor ions as a function of height in the corona as indicator of slow or fast wind

(Antonucci et al., 2006).


https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/L_SMALL_MRES_MCAD_Connection_Mosaic



Remarks:

• This is a SPICE-led observation. METIS cannot contribute to the study of the abundance,

but it can give the context of the coronal configuration before the limb off-pointing or farther

out.

• After METIS observations, we need to off-point to the limb in order to observe either a limb

active region and/or along the boundary of a streamer.

• SPICE will observe with a narrow slit in order to enhance the resolution for properly

measuring the abundance. A typical duration would be of 3.2 h.

• In the case of streamers, a bigger height can be studied from farther out, it is better to

observe from a distance >0.55 AU (for METIS support). In the case of active regions, the

perihelion is preferred for higher resolution.

• It might be difficult to address this objective at the maximum of solar activity.

• This observation should be tested at least once to see if it works and depending on the

success, re-plan it in the future.

The SOOP R_SMALL_HRES_LCAD_Composition_vs_Height is specifically designed for (and

currently only serves) this objective.

1.1.2.4 Study of density fluctuations in the extended corona as a function of the outflow

velocity of the solar wind while evolving in the heliosphere


Study of density fluctuations in the extended corona as a function of the outflow velocity of the

solar wind while evolving in the heliosphere (Telloni et al., 2009).

Remarks:

• A complete and detailed study of coronal density fluctuations requires a long temporal

baseline (more than one day).

• EUI/FSI 174 would need deep exposures in order to observe up to METIS FoV with good

SNR.

• METIS should observe with the FLUCTS mode for about 1 hour and then turn to MAGTOP

(with a cadence of 5-20 minutes). The current SOOP is defined with cycles of 8 hours, but

multiple cycles should be planned in order to observe for at least one day. Another

possibility would be to observe 8 hours per window for 3 consecutive windows, but this

might not meet the minimum observation duration for the lower frequencies.

• This objective should better be addressed at perihelion for a better METIS SNR for the

FLUCTS mode combined to a lower cadence during the rest of the window.

• If SPICE participates, we will need to do limb pointing for a certain period of time and then

return to disc center for METIS.

/solarorbiterwiki/display/SOSP/R_SMALL_HRES_LCAD_Composition_vs_Height



• PSP in quadrature would be an asset for observing the fluctuations in situ.

The SOOP R_FULL_HRES_HCAD_Density_Fluctuations is specifically designed for (and

currently only serves) this objective.

1.1.2.5 Structure and evolution of streamers


The extended observation at the limb of the same portion of the corona is essential to address the

study of streamer structure, evolution, and dynamics. This will focus either on the quiescent, slowly

evolving streamer belt or on the rapidly varying active streamers. The investigation of the

interrelation of solar wind acceleration region and magnetic topology of the flux tube guiding the

expansion necessitates as well long-term observations. The opportunity to freeze a streamer section

at the limb offers, in addition, the possibility of increasing dramatically the statistics, being, in this

case, only limited by the intrinsic evolution of the structure and not by solar rotation. A significant

increase in statistics is then coupled with the possibility of observing at high spatial resolution; this

allows us to resolve the fine structure and relevant dynamics in the slow wind coronal source

regions. This kind of studies requires the simultaneous knowledge of the electron density and

morphology of the corona, by observations in visible light, and of the radial flow velocities,

obtained with better accuracy and detail from images of the ultraviolet H I and He II emission, by

Doppler dimming techniques.

High-latitude observations will greatly impact on the study of large-scale structures of the solar

corona. If this phase occurs during solar minimum, from high latitudes the Orbiter views the

streamer belt running at low-latitude, or close to the equator, as an approximately continuous

annular structure around the solar disk. Furthermore, the out-of-ecliptic vantage point is indeed

providing another means to observe the corona not affected by solar rotation. When the morphology

is relatively simple and the relevant coronal features are at low-latitude, or close to the equator, their

intrinsic evolution can indeed be easily separated from rotational effects if viewed from high-

latitudes.

Another interesting aspect concerns the possibility to study the dynamics of coronal expansion all

around the streamer belt. During solar minimum, slow solar wind studies would be privileged, since

their supposed low-latitude and equatorial sources are predominant on the plane of the sky. It would

then be possible to (i) assess the contribution to the slow wind of sporadic reconnection events, such

as the coronal blowouts, and (ii) evaluate the total mass and magnetic flux injection into the

heliosphere all along the streamer belt.

Therefore, we have to:

• search for evidence of pseudo-streamers in the solar wind,

• study the detailed structure of the pseudo-streamers and

• study if slow wind is originating from pseudo-streamers.

Remarks:

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/R_FULL_HRES_HCAD_Density_Fluctuations



• Noci & Gavryuseva (2007) observed a quiescent streamer at solar minimum with UVCS for

4 consecutive days and determined the velocity pattern in it. Several days should be enough,

but several hours of observations could also give us useful insight.

• We should not observe at perihelion since we need METIS and a rather big FoV.

• A declining phase of the solar activity is preferred for a fully developed pseudo-streamer.

• Quadrature with Earth for combining with Solar Orbiter RS observations with L1 in situ, or

Earth coronagraphs observations with Solar Orbiter in situ.

This objective is addressed by the SOOP L_FULL_HRES_HCAD_Coronal_Dynamics, which is

aimed at observing structures in the outer corona and linking them to the heliosphere observed in

situ. METIS and SPICE are leading this SOOP, while in situ payload provides continuous

observations. Synoptic support from other remote sensing instruments is provided. Disk center

pointing is preferred.

1.1.2.6 Disentangle the spatial and temporal variability of the solar wind

Disentangle the spatial and temporal variability of the slow wind: comparison of the variability of

the slow wind on a number of orbits, with varying relative spacecraft-Sun longitudinal speeds at the

same distance. This objective is part of the general characterization of the solar wind and can be

addressed by the in-situ instruments. No specific operational planning is needed. However, it will

benefit from joint observations with PSP and L1, especially during radial alignments.

This objective is covered by the general in situ SOOP I_DEFAULT.

(To be moved to an overarching section (to be created) of characterization of the solar wind

together with other general in situ objectives).

1.1.2.7 Trace streamer blobs and other structures through the outer corona and the

heliosphere


• Trace streamer blobs ("helmet streamer plasmoids") and other structures through the outer

corona and the heliosphere (Sheeley and Rouillard, 2010).

• Study periodic density structures (Viall and Vourlidas, 2015) in the low corona and for

different times in the solar cycle.

The helmet streamer plasmoids are:

• observed in white light images as the continual, episodic releases of plasma from the tips of

helmet streamers

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• also, observed in situ at 1 AU and tracked from the upper corona to 1 AU (swept up and

compressed by the fast solar wind from low latitude coronal holes) (Sheeley and Rouillard,

2010).

• small interplanetary transients (size of 0.05-0.12 AU, magnetic field rotation between 2 and

13 hours) can be traced back to streamer events, but also to CMEs (Rouillard et al., 2011)

• flux ropes (Sheeley et al., 2009; Rouillard et al., 2011)

• exclusively plasma sheet phenomena (Wang et al., 2000)

• rate of ~4/day or approximately every 6 hours

• large in scale, initial sizes about 1 Rs in the radial direction and 0.1 Rs in the transverse

direction (Sheeley et al., 1997)

• originate at about 3-4 Rs (Sheeley et al., 1997)

• speeds: 150 km/s at 5 Rs, 300 km/s at 25 Rs

• thought to be released through either interchange reconnection, and/or complete

disconnection, and in either case, the reconnection takes place at high altitudes (Wang et al.

2000; Zurbuchen 2001; Crooker et al. 2004; Suess et al. 2009).

Other periodic density structures are:

• often not flux ropes (Viall et al., 2009)

• observed in 70%-80% of the slow solar wind and in much of the ecliptic fast wind (Viall et

al., 2008)

• rates of several minutes, characteristic timescale of 90 minutes, with a range of 65-100

minutes (Viall and Vourlidas, 2015)

• small scales (70-3000Mm)

• largest scale structures often contain smaller ones embedded within them (Viall et al., 2010)

• always associated with streamers and the HCS (Viall and Vourlidas, 2015)

• formed at or below 2.5 Rs, no evidence of formation above these heights (Viall and

Vourlidas, 2015).

Remarks:

• In general, not easy to distinguish radial vs transverse scales of structures. Joint observations

with PSP will be very important.

• At least one perihelion would be required for this objective for looking at the birth of the

structures. METIS would better observe from 0.3-0.35 during a zooming-out (North)

window.

• Low-latitude preferred.

• Minimum activity preferred, but the maximum could also be interesting if we could manage

to disentangle other effects.

• For linkage science, we would need Earth in quadrature and either observe in situ on Solar

Orbiter what Earth sees or observe in situ at 1 AU what SoloHI sees. For SoloHI to be

useful, we need a GSE –Y quadrature.

• The SoloHI turbulence mode should be used at perihelion for looking at the formation,

acceleration, shape, and evolution of the fastest periodic structures. For making sure that we



will see at least one blob at its birth (frequency of 1 every 6 hours), the turbulence mode (20

sec) should be used for approximately 12 hours (or 1 day).

This objective is addressed by the SOOP L_FULL_HRES_HCAD_Coronal_Dynamics, which is

aimed at observing structures in the outer corona and linking them to the heliosphere observed in

situ. METIS and SPICE are leading this SOOP, while in situ payload provides continuous

observations. Synoptic support from other remote sensing instruments is provided. Disk center

pointing is preferred.

Other SOOPs are partially addressing the connectivity aspects of the objective (but not the

formation and evolution of the fastest structures). These are

L_SMALL_MRES_MCAD_Connection_Mosaic,

L_SMALL_HRES_HCAD_SlowWindConnection, L_BOTH_MRES_MCAD_Farside_Connection.

1.1.2.8 Determine the velocity, acceleration profile and the mass of the transient slow wind

flows

This objective is part of the general characterization of the solar wind and can be addressed by the

in situ instruments. No specific operational planning is needed. However, it will benefit from joint

observations with PSP and L1, especially during radial alignments.

This objective is covered by the general in situ SOOP I_DEFAULT.

(To be moved to an overarching section (to be created) of characterization of the solar wind

together with other general in situ objectives).

1.1.3 Source regions of the heliospheric magnetic field

1.1.3.1 Full characterization of the photospheric magnetic fields and fine structures


• Full characterization of photospheric magnetic fields: the magnetic field at the photosphere

can be determined quantitatively by recording the full polarization state of light in

appropriate spectral lines. Such measurements can be inverted to provide the full magnetic

vector at the photosphere and the LOS component of the plasma flow velocity. These

measurements will allow the emergence of magnetic flux to be determined, as well as its

redistribution through its interaction with convection. From surface maps of the magnetic

vector, it is possible to extrapolate the field into the Sun’s upper atmosphere, where its

evolution gives rise to numerous dynamic and energetic phenomena. Time series of the

velocity maps of the photosphere also allow a reconstruction of the subsurface structure

using the techniques of helioseismology. Observations of the photospheric fields are, thus,

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https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/L_BOTH_MRES_MCAD_Farside_Connection



essential for studying both the generation and atmospheric evolution of solar magnetic

fields.

• Probe fine-scale structure of the solar photosphere, including waves and the emergence,

evolution, and dynamics of magnetic flux. This needs PHI observations at high spatial

resolution.

• Uncover the effects of waves and magnetic field changes on the upper atmospheric layers

(e.g. production of transient events (Galsgaard et al., 2005)).

• Study of the magnetic fields at the poles of the Sun.

o The polar fields are responsible for the polar coronal holes and largely drive the fast

solar wind, but are poorly known.

o Polar plumes are bright structures reaching far into the corona, which appear to

harbor gas moving slowly, compared to the fast solar wind in the interplume regions.

Observations of the magnetic field at their footpoints from high latitude will be

crucial for an understanding of their origin. The aim is to provide sufficiently

accurate and detailed magnetic maps, unhampered by the massive foreshortening that

current observations suffer from, to allow high-quality extrapolations of the field.

• Origin of the spicules and other chromospheric features (De Pontieu et al., 2004). Polar

spicules (Johannesson and Zirin, 1996). Spicules are a prominent chromospheric

phenomenon, cool and dense fibrils that intermittently connect the photosphere with the hot

and rarefied corona. They are short-lived (5-10 min), narrow (diameters less than 500 km)

and display upward motions with speeds up to 20 km/s. If this is really mass motion, then

the mass flux in spicules is 100 times larger than that of the solar wind. It has been shown

that photospheric p-modes, which are evanescent in the field-free photosphere and

chromosphere, can indeed propagate into and through the chromosphere if they are guided

by inclined magnetic field lines (De Pontieu et al. 2004). Due to the steep vertical density

gradient, the oscillations develop into shocks which may result in significant excursions of

the top of the chromosphere, i.e. cause spicules. According to De Pontieu et al. (2004), the

crucial ingredients for spicule formation are the photospheric velocities, the temperature

stratification and the inclination of the magnetic field lines. However, it is well known that

polar spicules are larger than ordinary spicules (Johannesson & Zirin 1996). They may,

therefore, differ in cause. Viewing them from out of the ecliptic will advantageously reveal

the underlying differences.

Remarks:

• Except for PHI/HRT (1 min cadence), EUI/HRI should observe at very high cadence (1-30s,

possible for approximately up to 20 minutes), maybe interleaved with burst modes with high

cadence (0.1 s, during 1 to a few seconds). The FoVs can be released if need be for telemetry

reasons (for instance down to ¼ of the PHI/HRT FoV).

• SPICE could participate for a dynamics study with many rasters and a similar FoV.

Alternatively, it could operate at a sit-and-stare mode.

• No other instruments are required.



• Due to its demanding nature in terms of telemetry, the relevant SOOP cannot be planned for

long durations. It should be repeated whenever the telemetry budget configuration is

favorable.

This objective is covered by the SOOP

R_SMALL_HRES_HCAD_Photospheric_Dynamics_Structure run for different targets at the quiet

Sun (observed from the perihelion) and the coronal holes (observed from high-latitude windows).

The SOOP L_SMALL_HRES_HCAD_SlowWindConnection will also provide valuable

information on photospheric structures but at lower spatial resolution.

1.1.3.2 How does the Sun’s magnetic field link into space?

For any plasma to leave the Sun and reach space, it must travel along the magnetic field. Any study

of how the Sun exerts its influence in space must, therefore, address the connection of the solar

magnetic field far into the solar system.

The large scale structure of the interplanetary magnetic field (IMF) is well known (e.g. Mariani and

Neubauer, 1990; Smith, 2008): the Sun's rotation winds the field into the Parker spiral; compression

and rarefaction of the plasma in co-rotating interaction regions (CIRs) produces increases and

decreases of the field strength; the polarity of the solar source field is reflected in that of the IMF;

and the field is pervaded by waves and turbulence over a wide range of scales. Over the Sun's 22

years magnetic cycle, the IMF reflects the changing character of the solar field, from approximate

dipole to a much more complex, multi-pole structure.

However, the mapping between solar and interplanetary fields is only known on relatively large

scales and in a crude manner. Observations of the Sun's surface photospheric magnetic field

combined with coronal observations or MHD models of the corona, make it possible to estimate the

mapping between the lower corona and the “source surface” at several solar radii, but many simple

questions remain about how the Sun's magnetic field opens into space (e.g. Antiochos et al., 2007),

particularly with regard to the emergence of new coronal holes and the long-range connectivity of

active regions, as well as how the IMF disconnects from the Sun. Distant observations by Ulysses

over the Sun's poles have helped to constrain such mappings (e.g. Hoeksema, 1995; Forsyth et al.,

1997) but Solar Orbiter, being much closer to the Sun and hence eliminating many of the

uncertainties caused by local stream-stream interactions, will dramatically improve the precision

with which this can be constrained.

Beyond the source surface, MHD models must be used. These models will be greatly constrained by

Solar Orbiter magnetic field data, with the important consequence of improving the systematic

prediction ability of such models throughout the heliosphere. This connection is also essential for

many elements of the Solar Orbiter science objectives of linking solar and interplanetary

phenomena.

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1.1.3.2.1 How does the Sun’s magnetic field change over time?


The process by which the Sun's field reverses at solar maximum is highly complex: lack of

observations of the polar solar field greatly hampers our understanding of this process. Ulysses in

situ observations contributed to measurements of the previous cycle's reversal (e.g. Jones et al.,

2003). Solar Orbiter will image the polar field with PHI while simultaneously measuring the field in

space at a range of locations, making more precise measurements of the solar reversal and its effects

on the heliosphere. While the nature of the solar field reversal is part of Objective 4 (How does the

solar dynamo work and drive connections between the Sun and the heliosphere?), here we are

interested in how this affects the coronal and heliospheric magnetic fields.

We will measure how the polarity and large scale structure of the Sun's magnetic field in

interplanetary space varies close to the Sun as it moves from solar minimum towards maximum and

the global field reverses.

Even though part of this objective can be addressed with in situ only measurements, it would make

more sense to have access to full disk remote sensing observations in order to properly understand

how the in-situ changes are linked to the overall changes in the solar field. For this, we need long-

term observations at synoptic modes (for telemetry reasons), by targeting the full disk for both

photospheric and coronal fields. Since such (low resolution) observations already exist from Earth,

it would mostly be interesting to observe when Solar Orbiter is at the far side of the Sun or for

intermediate regions between an alignment with the Earth and the far side (regular spacing in

longitude). This should also be repeated for different latitudes. Since in situ MAG measurements are

key for this objective, good statistics during EMC Quiet periods are required.

This objective can be addressed by SOOP L_FULL_HRES_LCAD_MagnFieldConfig, which has

been mainly created for this goal.

1.1.3.2.2 How is the heliospheric current sheet (HCS) related to coronal structure?


The Heliospheric Current Sheet (HCS) is the interplanetary extension of the neutral line between the

two magnetic polarities in the corona and is a topological magnetic boundary. The HCS is also vital

in the motion of cosmic rays throughout the heliosphere: depending on the polarity of the solar

cycle, ions or electrons tend to migrate to low latitudes and along the HCS as they enter the solar

system.

The HCS is remarkably thin: just a few thousand km across (Zhou et al., 2005), but is surrounded by

the much thicker, denser Heliospheric Plasma Sheet (HPS) (Vourlidas and Riley, 2007). It is not

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clear how thin the HCS and HPS are close to the Sun, which could provide clues to their origin – the

HCS seems actually to thin with distance, for example. Both the HCS and HPS are also highly

variable: what is the origin of this variability? Reconnection appears to occur here (Gosling et al.,

2006): how frequent is this close to the Sun?

Data from a number of sources (cosmic rays: Simpson et al., 1996; the IMF: Luhmann et al., 1988;

geomagnetic activity: Mursala and Zieger, 2001; source surface models: Zhao et al., 2005) suggest

that the Sun's neutral line, and also the HCS, are persistently displaced Southward during solar

minimum. This is consistent with a solar dipolar field offset from the Sun's center, but why (or even

if) this should occur is unknown.

We, therefore, have to:

• Determine the local tilt and latitudinal extent of the HCS as a function of solar distance and

relation to the time-varying source surface neutral line. Asymmetry of the HCS. Determine

whether the HCS is offset from the equator in near-Sun space.

• Study the relation of the Heliospheric Plasma Sheet (HPS) to the HCS. (Vourlidas and Riley,

2007).

• Measure the variability of the HCS and HPS in time and space, with the goal of determining

the solar or local origin of the variations. Reconnection (Gosling et al., 2006)?

• Determine the link between heliospheric plasma sheet and coronal streamers (Bavassano et

al., 1997).

Remarks:

• As for objective 1.1.3.2.1 How does the Sun's magnetic field change over time?, in situ

observations are key during EMC Quiet periods, but full disk synoptic remote sensing

observations will also be important for a better understanding.

• This objective can be addressed by SOOP L_FULL_HRES_LCAD_MagnFieldConfig.

• It would be interesting to study the HCS crossings at the perihelia during the Cruise Phase.

• Observations during one solar rotation will not be enough since we will have gaps due to the

RS windows durations. High-latitude observations are also required.

1.1.3.2.3 How does the heliospheric magnetic field disconnect from the Sun?


At any given time, magnetic field lines stretch out from the Sun deep into interplanetary space,

carried by the solar wind. However, these fields must eventually disconnect from the Sun, resulting

in a complex, tangled magnetic topology in the heliosphere, with important consequences for

energetic particle propagation. It is essential to quantify the large-scale connectivity, but this is

remarkably difficult. The magnetic field polarity is an important indicator of the solar region of

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origin of a packet of solar wind, but cannot determine its connectivity alone. Signatures of

connectivity such as suprathermal electron streaming (e.g. Zurbuchen and Richardson, 2006) are

difficult to interpret. Indeed, (Owens and Crooker, 2007) suggests that electron scattering and fading

can result in signatures which look like disconnected field lines, but are actually due to interchange

reconnection. Only by simultaneously measuring the streaming electrons and magnetic field polarity

over a wide range of distances, and particularly close to the Sun, can we distinguish dis- and re-

connections from electron fade-out.

Evidence of reconnection at solar wind magnetic discontinuities (e.g. Gosling et al., 2007)

demonstrates that the connectivity of the IMF can change in space. How often does this occur close

to the Sun, where the solar wind is much more dynamic than at 1 AU? There is also evidence for

folds in the magnetic field, from cross-helicity (Balogh et al., 1999) and proton-alpha streaming data

(Yamauchi et al., 2004): what is their origin? Are they related to chromospheric reconnection

features such as jets (Shibata et al., 2007) or velocity shears (Landi et al., 2006)?

Remarks:

We have to search for local reconnection events using MAG and SWA in order to determine their

radial distribution and significance for the connectivity of the solar wind.

We will use signatures of the varying connectivity of the solar wind (e.g. suprathermal electrons),

combined with the magnetic field orientation and other measures of source polarity (e.g.

alpha/proton streaming and the normalized cross helicity) to search for bends, folds and small scale

polarity reversals. In this way, we will determine the small-scale polarity structure within coronal

hole flows and its relation to the global field.

This is an in situ objective and can be addressed with the I_DEFAULT SOOP as well as during the

connectivity operating plans for the fast and slow winds,

e.g. L_SMALL_HRES_HCAD_Fast_Wind, L_SMALL_HRES_HCAD_SlowWindConnection, L_

SMALL_MRES_MCAD_Connection_Mosaic.

1.1.3.3 What is the distribution of the open magnetic flux?


• Understand what determines the amount of open flux from the Sun, how open field lines are

distributed at the solar surface at any given time, and how these open field lines reconnect

and change their connection across the solar surface in time.

• Origin of the open magnetic flux from:

▪

▪ Coronal holes.

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▪ Active regions.

▪ Quiet Sun.

▪ Loops of varying heights.

▪ Tips and legs of the helmet streamer.

▪ Polar plumes.

▪ Other solar structures.

Remarks:

• We need MAG measurements for the magnetic field polarity and strength, over a wide range

of distances and latitudes, combined with signatures of connectivity such as suprathermal

electrons (SWA/EAS) to determine the amount of open magnetic flux in the heliosphere

from solar minimum to maximum.

• METIS can provide electron density and outflow velocity maps with a 20-30 min cadence.

• EUI can support with FSI for complementing observations in the lower lying corona, at low

cadence as METIS.

• PHI can provide low-resolution synoptic context.

• As for all goals of the category 1.1.3 Source regions of the heliospheric magnetic field, far

side observations will be most innovative, as well as high-latitude. EMC Quiet periods are

required for the MAG measurements.

Part of this objective can be addressed with the I_DEFAULT SOOP as well as during the

connectivity operating plans for the fast and slow winds,

e.g. L_SMALL_HRES_HCAD_Fast_Wind, L_SMALL_HRES_HCAD_SlowWindConnection, L_

SMALL_MRES_MCAD_Connection_Mosaic.

However up to now, we have only included it in R_SMALL_MRES_MCAD_AR_LongTerm, but

this would only partially address it.

1.1.4 Transverse themes (reconnection)

This section will be modified or completely removed by integrating the different parts in other

objectives. For this document, I keep it empty for now.

3.2.1.2 1.2 What mechanisms heat the corona and heat and accelerate the solar wind?

Present state of knowledge:






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Despite more than a half-century of study, the basic physical processes responsible for heating the

million-degree corona and accelerating the solar wind are still not known. Identification of these

processes is important for understanding the origins and impacts of space weather and to make

progress in fundamental stellar astrophysics.

Ultimately, the problem of solar wind acceleration is a question of the transfer, storage, and

dissipation of the abundant energy present in the solar convective flows. The key question is to

establish how a small fraction of that energy is transformed into magnetic and thermal energy above

the photosphere. Both emerging magnetic flux and the constant convective shaking and tangling of

magnetic field lines already threading the corona contribute to the processing of the energy in what

is an extremely structured, highly dynamic region of the solar atmosphere, the route to dissipation

involving cascading turbulence, current sheet collapse and reconnection, shocks, high-frequency

waves, and wave-particle interactions. The advent of high-cadence high-resolution observations has

demonstrated the extremely complex phenomenology of the energy flux in the lower atmosphere,

including many types of transient events discovered and classified by Yohkoh, SOHO, TRACE,

RHESSI, Hinode and, most recently, SDO.

Energy deposited in the corona is lost in the form of conduction, radiation (negligible in coronal

holes), gravitational enthalpy, and kinetic energy fluxes into the accelerating solar wind plasma.

Transition region pressures, coronal densities and temperatures, and the asymptotic solar wind speed

are sensitive functions of the mode and location of energy deposition. The mass flux is not,

however, as it depends only on the amplitude of the energy flux (Hansteen and Leer 1995). A

relatively constant coronal energy flux therefore explains the small variations in mass flux between

slow and fast solar wind found by Ulysses during its first two orbits, although the dramatic decrease

in mass flux over the last cycle points also to a decreased efficiency of coronal heating and therefore

to its dependence on the solar magnetic field (McComas et al. 2008; Schwadron and McComas

2008).

One of the fundamental experimental facts that has been difficult to account for theoretically is that

the fast solar wind originates in regions where the electron temperature and densities are low, while

the slow solar wind comes from hotter regions of the corona. The anticorrelation of solar wind

speed with electron temperature is confirmed by the anti-correlation between wind speed and

‘freezing in’ temperature of the different ionization states of heavy ions in the solar wind (Geiss et

al. 1995) and implies that the electron pressure gradient does not play a major role in the

acceleration of the fast wind. On the other hand, the speed of the solar wind is positively correlated

with the in-situ proton temperature, and the fastest and least collisionally coupled wind streams also

contain the largest distribution function anisotropies. Observations of the very high temperatures

and anisotropies of coronal heavy ions suggest that other processes such as magnetic mirror and

wave-particle interactions should also contribute strongly to the expansion of the fast wind (Li et al.

1998; Kohl et al. 1997, 1998, 2006; Dodero et al. 1998). In particular, either the direct generation of

high-frequency waves close to the cyclotron resonance of ions or the turbulent cascade of energy to

those frequencies should play an important role.

Theoretical attempts to develop self-consistent models of fast solar wind acceleration have followed

two somewhat different paths. First, there are models in which the convection-driven jostling of

magnetic flux tubes in the photosphere drives wave-like fluctuations that propagate up into the

extended corona. The waves partially reflect back toward the Sun, develop into strong turbulence,

and/or dissipate over a range of heights. These models also tend to attribute the differences between



the fast and slow solar wind not to any major differences in the lower boundary conditions, but to

the varying expansion factor of magnetic field lines in different areas of coronal holes (Cranmer et

al. 2007).

In the second class of models, the interchange reconnection models, the energy flux usually results

from magnetic reconnection between closed, loop-like magnetic flux systems (which are in the

process of emerging, fragmenting, and being otherwise jostled by convection) and the open flux

tubes that connect to the solar wind. Here the differences between fast and slow solar wind result

from qualitatively different rates of flux emergence, reconnection, and coronal heating in different

regions on the Sun (Axford and McKenzie 1992; Fisk et al. 1999; Schwadron and McComas, 2003).

It has been difficult to evaluate competing models of fast wind acceleration and to assess

observationally the relative contributions of locally emerging magnetic fields and waves to the heat

input and pressure required to accelerate the wind largely because of the absence of measurements

of the solar wind close to the Sun where they can be mapped with sufficient precision to a solar

source region.

How Solar Orbiter will address this question:

Solar Orbiter’s combination of high-resolution measurements of the photospheric magnetic field

together with images and spectra at unprecedented spatial resolution will make it possible to identify

plasma processes such as reconnection/shock formation and wave dissipation in rapidly varying

surface features, observe Doppler shifts of the generated upflows, and determine

compositional signatures. Whatever the scale, magnetic reconnection leads to particle dissipative

heating and acceleration and wave generation, which have the net effect of a local kinetic energy

increase in the lower solar atmosphere that can be revealed through high-resolution extreme

ultraviolet (EUV) imaging and spectroscopy. Wave propagation will be traced from the source site

to the region of dissipation through observations of EUV-line broadening and Doppler shifts.

Global maps of the H outflow velocity, obtained by applying the Doppler dimming technique to the

resonantly scattered component of the most intense emission line of the outer corona (H I 121.6),

will provide the contours of the maximum coronal expansion velocity gradient for the major

component of the solar wind, and the role of high-frequency cyclotron waves will be

comprehensively assessed by measuring spectroscopically the particle velocity distribution across

the field and determining the height where the maximum gradient of outflow velocity occurs

(Telloni et al. 2007).

Solar Orbiter’s heliospheric imager will measure the velocity, acceleration, and mass density of

structures in the accelerating wind, allowing precise comparison with the different acceleration

profiles of turbulence- driven and interchange reconnection-driven solar wind models.

As it is performing imaging and spectroscopic observations of the corona and photosphere, Solar

Orbiter will simultaneously measure in situ the properties of the solar wind emanating from the

source regions. The in-situ instrumentation will determine all of the properties predicted by solar

wind acceleration models: speed, mass flux, composition, charge states, and wave amplitudes.

Moving relatively slowly over the solar surface near perihelion, Solar Orbiter will measure how

properties of the solar wind vary depending on the changing properties of its source region, as a

function of both space and time, distinguishing between competing models of solar wind

generation.



1.2.1 What mechanisms heat the corona?

The detailed objectives for understanding the heating of the corona include, but are not limited to,

the following:

1.2.1.1 Energy flux in the lower atmosphere


• Understand dissipation involving cascading turbulence, current sheet collapse and

reconnection, shocks, high-frequency waves and wave-particle interactions. Difference

between resonant and stochastic?

• Understand oscillations in small-scale flux tubes (Jess et al., 2009; Martínez González, 2011;

Stangalini et al., 2013; Requerey et al., 2015).

The scientific aim is to characterize the properties of waves in the photosphere and their coupling

with the upper atmosphere, chromosphere, and corona.

Waves are one clear mechanism for transferring energy from the photosphere to the chromosphere

and corona. Measuring the properties of the waves requires, in part, a determination of the velocity

field. The line-of-sight velocity component can be determined at different heights in the atmosphere

by observing Doppler shifts in different spectral lines. From the Earth’s vantage point we have high-

resolution ground-based, balloon-borne, and satellite instruments. Determining the horizontal

velocity has previously relied on using correlation tracking of intensity variations and rely on the

questionable assumption that the changes in the location of the brightness fluctuations reflect the

actual velocity. The orbit and capability to measure Doppler velocities, in conjunction with existing

and upcoming ground-based or near-earth observatories, offers the unique chance to directly

measure two components of the velocity field using the Doppler effect.

Remarks:

High-resolution co-temporal measurements including Doppler velocity maps from Solar Orbiter as

well as ground and NEOs are required. In particular, the ground-based and NEOs should include

high-resolution Doppler images in the same line (with a higher cadence than that of SO), as well as

lines sampling different heights of the atmosphere. Co-observation with IRIS would be desirable.

During the observing period, the Earth-Sun-SO angle should be between 30◦ and 60◦ – a range

which represents a compromise between determining the two components of the velocity field and

allowing magnetic features which can act as wave guides to be partially resolved.

For ease of understanding the connection between the different heights, the observations would best

be performed at the center of the disk as observed from the Earth (where observations over different

wavelengths are possible). Because also the achievable cadence will be higher on ground than

with PHI, it is preferable to select targets which are closer to the disk center as seen from Earth and

at higher heliocentric angles as seen from Solar Orbiter. The highest possible cadence is desirable,

and a shorter time series (of down to 30 minutes of Solar Orbiter observations) would still allow the



scientific objectives to be met. (The ground-based and NEO should be made for a period of 90

minutes centered on the 30 minute SO observations). However, in order to guarantee reliable

conditions (seeing) at the coordinating ground-based observing facility (e.g. DKIST) a continuous

high-cadence observation period of several hours is required.

High-resolution context magnetic maps from SO immediately before and after the 30 minute

observing window are required to provide context and aid co-alignment.

A second observational campaign of an area 45◦ from disk center, with an Earth-Sun-SO angle of

90◦, would be desirable.

This objective can be partly covered by the high-resolution high-cadence SOOP that studies the fine

scale of the photosphere R_SMALL_HRES_HCAD_RSburst and the

SOOP R_SMALL_HRES_HCAD_WaveStereoscopy, which has been specifically designed for

studying the properties of the waves in the photosphere.

The SOOPs should be run for a bright source (e.g. active region but also quiet sun) at the perihelion

for a duration from 20 minutes to several hours. It would be preferable to address this objective

earlier in the mission, because of the Lyα degradation with time.

1.2.1.2 Energy and mass flux in the corona. Loss in form of conduction, radiation, gravitational

enthalpy, and kinetic energy fluxes into the accelerating solar wind plasma.


Energy and mass flux in the corona. Loss in form of conduction, radiation, gravitational energy,

enthalpy, and kinetic energy fluxes into the accelerating solar wind plasma.

This goal needs better definition and detailed description.

1.2.1.3 Contribution of flare-like events on all scales


• Statistical study of more than 70k flare-like events on all scales that are expected to be

seen during the NMP for better defining the flares’ distribution that possibly contains

enough energy to heat the corona (Parker 1983; Hudson, 1991).

• Reduce the lower energy limit of detectable flares.

• Understand the mechanisms responsible for producing nanoflares (field line shear, braiding,

coronal tectonics, (Priest, 2003)).

• Search for small 3He-rich events close to the perihelion.

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1.2.1.4 Observe and explore flare-like ‘heating events’ from the quiet corona

• Observe and explore flare-like ‘heating events’ from the quiet corona (Krucker et al., 1997).

• Possibility to observe or not hard X-ray counterparts not yet detected (Hannah et al., 2007).

1.2.1.5 Determine whether coronal heating is spatially localized or uniform, and time steady or

transient or impulsive for a wide range of magnetic loops with different spatial scales. Observe

coronal nanoflares in active region moss (Testa et al., 2013).

1.2.1.6 Resolve the geometry of fine elemental loop strands (Aschwanden et al., 2002; Reale et al.,

2007) and define the possible multi-temperature nature of such structures and their transversal

expansion in the corona, giving insight to the level of the tangling of the magnetic flux tubes (Lopez

Fuentes et al., 2006).

1.2.1.7 Detect and characterize waves in closed and open structures (e.g. De Pontieu et al., 2007),

looking for signatures of their dissipation in the transition region and low corona, with a

corresponding evaluation of the energy released to heat the solar plasma.

1.2.1.8 Investigate the role of small scale magnetic flux emergence in energizing the above laying

layers (e.g. Galsgaard et al., 2007).

1.2.1.9 Multiple-temperature diagnostics of flaring coronal loops (Battaglia and Kontar, 2013).

1.2.1.10 Heating in flaring loops vs in active regions.

1.2.2 What mechanisms heat and accelerate the solar wind?

1.2.2.1 Determine where energy is deposited in the solar wind (Cranmer et al. 1997; Telloni et al.

2007).

1.2.2.2 What drives the evolution of the solar wind distribution?


• High time resolution variability in the magnetic field associated with particle distributions in

different plasma parameter regimes (low and high plasma β, fast and slow wind) in order to

quantify the links between particle distribution evolution and wave-particle interactions

(Kasper et al., 2002; Matteini et al., 2007; Maksimovic et al., 2005).

• Role of the electron heat flux. Characterize the non-thermal character of the electron

distributions at perihelion and their evolution with heliocentric distance (Stverak et al., 2009;

Maksimovic et al., 2005).



• Generation of non-thermal ion distributions (beams, drifting heavy ions, hot heavy ions,

proton ‘strahl’).

• Generation of non-thermal electron distributions (beams, drifting ‘strahl’).

• What role do local evolution (turbulence, shears), global evolution (expansion), and

collisions have in determining the properties of the proton distributions?

• Sub-Debye length electric fields (Randol & Christian, 2014), measure electric fields and

suprathermal proton tails.

1.2.2.3 What are the origins of waves, turbulence and small scale structures?

• Quantify the undisturbed waves and relate the wave power and other characteristics to the

source regions (by measuring photospheric motion in the regions from which the plasma

originated) (Matthaeus and Goldstein, 1986; Bruno and Carbone, 2005; Matthaeus et al.,

2007; Bello González et al., 2010).

• Identify and characterize the waves associated with the plasma instabilities that isotropize

and heat the solar wind (Hellinger et al., 2006; Matteini et al., 2010).

• Resonant absorption and emission by thermal particle distributions: role of the high-

frequency cyclotron waves.

• Ion energization processes in the solar wind (study of the electric fluctuations near the ion

cyclotron frequencies).

• Ion cyclotron resonance damping of the high-frequency part of the Alfvén spectrum (e.g.

Cranmer, 2002).

• Solve the problem of the mode conversion from Langmuir to electromagnetic waves (Bale et

al., 1998; Kellogg et al. 1999; Farrell et al., 2004; Ergun et al., 2008). Characterize the

energy balance between electron beams, Langmuir waves and electromagnetic radio waves

at several heliocentric distances.

• How do variations and structure in the solar wind affect low-frequency radio wave

propagation?

• Small scale structures such as solitons (Rees et al., 2006), mirror modes (Stevens and

Kasper, 2007), and draped fields and - in the case of dust trail signatures (Jones et al., 2003)

- confirm or refute their correlation with predicted trails.

• Study inbound waves in the corona (Verdini et al., 2009; DeForest et al., 2014).

1.2.2.4 Identify and characterize the solar wind reconnection physics in current sheets with

thickness down to the ion scales and smaller. Electron and ion physics near the current sheet.

Magnetic islands formation. Compare microphysics of solar wind reconnection with

magnetospheric reconnection.

1.2.2.5 Magnetic reconnection in the chromosphere, the transition region and the corona, driven by

magnetic field evolutionary processes. Explore reconnection signatures such as brightenings

(Harrison et al., 1997), flows and jets (Innes et al., 1997) and plasma evaporation (Klimchuk, 2006).

Reconnection between small closed loops and open structures supplying energy to the nascent solar

wind in the form of waves and turbulent flows (Tu et al., 2005).



1.2.2.6 Study fast plasma flows from the edges of solar active regions discovered with Hinode/EIS

(Harra et al., 2008), which are driven by pressure gradient between reconnecting magnetic loops

(Baker et al., 2009; Del Zanna et al., 2011) and produce intermediate-speed solar wind streams (van

Driel-Gesztelyi et al., 2012). Study active region expansion (Uchida et al., 1992).

1.2.2.7 Study the correlation degree between velocity and magnetic field fluctuations in the

interplanetary space: may help in disentangling the source regions of the slow solar wind.

3.2.1.3 1.3 What are the sources of solar wind turbulence and how does it evolve?


The solar wind is filled with turbulence and instabilities. At large scales, the fast solar wind is

dominated by anti-sunward propagating Alfvén waves thought to be generated by photospheric

motions. At smaller scales, these waves decay and generate an active turbulent cascade, with a

spectrum similar to the Kolmogorov hydrodynamic scaling of f^-5/3. In the slow solar wind,

turbulence does not have a dominant Alfvénic component, and it is fully developed over all

measured scales. There is strong evidence that the cascade to smaller scales is anisotropic, but it is

not known how the anisotropy is generated or driven (Horbury et al. 2008). What do the differences

between the turbulence observed in the fast wind and that observed in the slow wind reveal about

the sources of the turbulence and of the wind itself?

Little is known about what drives the evolution of solar wind turbulence. Slow-fast wind shears,

fine-scale structures, and gradients are all candidate mechanisms (Tu and Marsch 1990; Breech et

al. 2008). To determine how the plasma environment affects the dynamical evolution of solar wind

turbulence it is essential to measure the plasma and magnetic field fluctuations in the solar wind as

close to the Sun as possible, before the effects of mechanisms such as velocity shear become

significant, and then to observe how the turbulence evolves with heliocentric distance.

The dissipation of energy in a turbulent cascade contributes to the heating of the solar wind plasma.

However, while measurements of the properties of solar wind turbulence in near-Earth orbit largely

agree with observed heating rates (Smith et al. 2001; Marino et al. 2008), the details are

controversial and dependent on precise models of turbulent dynamics. In order to establish a full

energy budget for the solar wind, the heating rates as a function of distance and stream properties

must be determined, including turbulence levels before the cascade develops significantly.

The statistical analysis of the fluctuating fields also reveals pervasive fine-scale structure (e.g.,

discontinuities and pressure balanced structures). The origin of these structures is uncertain: are they

the remnant of complex coronal structuring in the form of strands of small-scale flux tubes advected

by the solar wind flow (Borovsky 2008; Bruno et al. 2001), or are they generated locally by

turbulent fluctuations?

At scales around the proton gyroradius and below, turbulent fluctuations interact directly with the

solar wind ions. The precise nature of the turbulent cascade below the proton gyroradius is poorly

understood and might even vary depending on local plasma conditions. Below the electron



gyroradius, conditions are even less certain and the partitioning of turbulent energy into electron or

ion heating is unknown at this time. In addition, solar wind expansion constantly drives distribution

functions toward kinetic instabilities, where fluctuations with characteristic signatures are generated

(e.g., Marsch 2006). What physical role do kinetic effects play with distance from the Sun? What

role do wave-particle interactions play in accelerating the fast solar wind? What contribution do

minor ions make to the turbulent energy density in near-Sun space?


Solar Orbiter will measure waves and turbulence in the solar corona and solar wind over a wide

range of latitudes and distances, including closer to the Sun than ever before, making it possible to

study turbulence before it is significantly affected by stream-stream interactions. By traveling over a

range of distances, the spacecraft will determine how the turbulence evolves and is driven as it is

carried anti-sunward by the solar wind flow.

Detailed in-situ data will make it possible to distinguish between competing theories of turbulent

dissipation and heating mechanisms in a range of plasma environments and are thus of critical

importance for advancing our understanding of coronal heating and of the role of turbulence in

stellar winds.

By entering near-corotation close to the Sun, Solar Orbiter will be able to distinguish between the

radial, longitudinal, and temporal scales of small-scale structures, determining whether they are the

signatures of embedded flux tubes or are generated by local turbulence.

Solar Orbiter’s magnetic and electric field measurements, combined with measurement of the full

distribution functions of the protons and electrons will fully characterize plasma turbulence over all

physically relevant time scales from very low frequencies to above the electron gyrofrequency.

Because Solar Orbiter is a three-axis stabilized spacecraft, it can continuously view the solar wind

beam with its proton instrument, measuring proton distributions at the gyroperiod and hence making

it possible directly to diagnose wave-particle interactions in ways that are not possible on spinning

spacecraft. By traveling closer to the Sun than ever before, it will measure wave-particle interactions

before the particle distributions have fully thermalized, studying the same processes that occur in

the corona. By measuring how the distributions and waves change with solar distance and between

solar wind streams with different plasma properties, Solar Orbiter will make it possible to determine

the relative effects of instabilities and turbulence in heating the plasma.

The solar wind is the only available plasma laboratory where detailed studies of

magnetohydrodynamic (MHD) turbulence can be carried out free from interference with spatial

boundaries and in the important domain of very large magnetic Reynolds numbers. A detailed

comparison between experimental in-situ data and theoretical concepts will provide a more solid

physical foundation for MHD turbulence theory, which will be of critical importance for

understanding the solar coronal heating mechanism and the role or turbulence in the solar wind.

1.3.1 Solar and local origin of Alfvénic fluctuations




• Outward propagating stochastic fluctuations, Alfvénic modes within the fast streams

(Belcher and Davis, 1971; Tu and Marsch, 1995).

o Excitement by magnetic activity and reconnection in the chromospheric network.

o Large-scale torsional Alfvén waves generated from magnetic restructuring related to

interchange reconnection (Lynch et al., 2014).

o Alfvén waves in the low corona caused by leakage of wave power from

chromospheric oscillations (Tomczyk et al., 2007).

o Presence or not of large-amplitude Alfvénic fluctuations with ⎟B⎜=const (Matteini

et al., 2014) in the inner heliosphere, implication for their stability and their role as a

source for turbulence.

• Identification of the drivers at the "outer scale" (large scale) end of the turbulent cascade.

• Inward propagating Alfvén waves (already active at 0.3 AU as observed by Helios):

o Local production by velocity shears (Roberts et al., 1992).

o Parametric decay mechanisms (Malara et al., 2000).

o Generation of inward component by expansion (relevant for polar wind where shear

effects are negligible) (Velli et al., 1989; Verdini et al., 2009).

o Nonlinear generation by compressive MHD cascade (Marsch & Mangeney, 1987).

o Carefully distinguish real local production of inward modes from compressive events

masking a generation process (Bavassano and Bruno, 1989).

• How does the behavior of MHD waves observed remotely in the core of coronal holes and at

the CH/QS boundaries match what is observed in situ in the solar wind that emanates from

these regions?

Remarks:

• Note that the McIntosh paper (McIntosh et al., Nature, 475, 477) observed waves in

structures seen above the limb (spicules) and in active region loops (on-disc). In coronal

holes seen on the disc, it will be probably easier to observe oscillating structures in Lyman

alpha (HRI) than in the other passbands. In Ly-alpha, there are dark, fibril-like structures

that are seen with a better contrast with respect to the background (optical thickness), while

in 304 or 174, the structures are optically thin and the contrast may be less favorable. But

using HRI and SPICE will probably require off-pointing, so we should implement two

different observational strategies: one with and one without off-pointing.

• For the objective of linking with in-situ observations, the "minimal" option should be

without SPICE or EUI/HRI, and no off-pointing required. EUI/FSI 174 and 304 at 30-s

cadence during 1-2 hours minimum (there is currently no program fitting this cadence with

FSI). Alternating both filters would probably require filter wheel motions incompatible with

EMC requirements. If so, FSI 304 may be preferred. Or 1 hour of FSI 304 followed by 1

hour of FSI 174. To save telemetry, it is possible to store only 1/4 of the FOV as long as it

covers the core and one boundary of the CH. The 30-s cadence is sufficient to observe waves

as in McIntosh et al., Nature, 475, 477. But the spatial resolution of FSI may not be

sufficient to resolve fine structures and their motions, especially for observations on the disc,

which are necessary as the objective is to measure the wind emanating from the RS

observations.



• "Optimal" option: with SPICE and HRI; off-pointing and mosaic may be necessary to

observe the source region of the wind that will be measured at the spacecraft within a few

days from the RS observations. The observed region must cover both the core of the CH and

at least one of the CH/QS boundaries.

• HRI 174 and HRI Ly-alpha at 30-s cadence during 1-2 hours (EUI Coronal hole mode (C)

can fit). Observations in coronal holes require deep exposures (at least in HRI 174), which

requires minimal jitter.

• FSI 304 at 30-s cadence.

• SPICE rasters for measurement of the line widths (to determine the non-thermal motions,

supposed to correspond to LOS-integrated effect of Alfvén waves), for at least transition

region and coronal lines (optionally: also, lines for abundance measurements, for which the

total intensity may be enough and not the full profile). Cadence is not an issue as we

measure the integrated effect in the line width (for observation of the wave motions in

Doppler shifts, sit-and-stare (i.e. no slit scan) may be preferable).

1.3.2 How is turbulent energy dissipated and how does turbulence evolve within the heliosphere?


• Measure the turbulence dissipation range and understand its scaling with heliocentric

distance and plasma properties (Alexandrova et al., 2007; Bruno & Trenchi, 2014; Chen et

al., 2014). Separate the analysis between fast and slow wind and in particular for different

plasma beta.

• Distinguish between various heating and dissipation mechanisms. Various kinetic processes:

Alfvénic or magneto sonic damping, kinetic Alfvén waves (see above), whistler dispersion,

Hall MHD dispersive cascade.

• Study the evolution of the intermittency of the magnetic and plasma quantities (Bruno et al.,

2003; 2014).

• Study the evolution of the effective magnetic Reynolds number (Bruno et al., 2015).

• Study the evolution of the MHD rugged invariants (magnetic helicity, cross-helicity, and

residual energy).

• Relate the localization of the spectral break between the fluid and the kinetic regimes to the

amplitude of the fluctuations at MHD scales: are the scales at which dissipation mechanisms

become important related to the energy contained in the inertial range? (Bruno and Trenchi,

2014).

• Explore the radial evolution of the compressible and incompressible third order moment

scaling within the inertial range, in order to gain information on the status of the turbulent

cascade (Sorriso-Valvo et al., 2007), including: the evaluation of the mean energy transfer

rate, the development of the cascade relative to the balance between the inward/outward

fluctuations, the dependence on local parameters.

• Understand the origin and the radial evolution of the low frequency 1/f spectrum in fast solar

wind and its implications on the characteristics of turbulence, intermittency and dissipation.

• Dissipation at the corona.



1.3.3 Plasma turbulence variability


Quantify the evolution of the turbulent spectrum with distance, heliomagnetic latitude and stream

structure in the inner heliosphere.

• Determine in what conditions the radial evolution is a matter of age of turbulence or in situ

heliospheric sources.

• Try to connect in situ spectrum with source spectrum at the Sun (together with sophisticated

models).

• Get insight into the observed correlation between speed, temperature and fluctuations

amplitude (Grappin et al., 1990) by comparing in situ properties and source region

properties.

1.3.4 Plasma turbulence anisotropy


Quantify the anisotropy of the turbulence in both slow and fast streams; determine its effects on

particles at a range of energies and relate it to the solar origin of the fluctuations (Horbury et al.,

2005; Dasso et al., 2005).

• Determine if different levels of anisotropy result from differences in the source region or

from slower/faster evolution in fast/slow streams (age of turbulence).

• Quantify the effect of expansion (radial evolution) on the 3D anisotropy of turbulence as

observed at 1AU (Narita et al., 2010) and predicted by MHD simulations (Dong et al.,

2014).

• Compare turbulence anisotropy to measured diffusion coefficients of energetic particles

during solar events (Horbury and Balogh, 2001; Bieber et al., 1996).





3.2.2 Objective 2: How do solar transients drive heliospheric variability?

The dynamic Sun exhibits many forms of transient phenomena, such as flares, CMEs, eruptive

prominences, and shock waves. Many directly affect the structure and dynamics of the outflowing

solar wind and thereby also eventually affect Earth’s magnetosphere and upper atmosphere, with

significant consequences for society through hazards to, for example, space-based technology

systems and surface power systems. Understanding these impacts, with the ultimate aim of

predicting them, has received much attention during the past decade and a half under the banner of

‘space weather.’ However, many fundamental questions remain about the physics underpinning

these phenomena and their origins, and these questions must be answered before we can realistically

expect to be able to predict the occurrence of solar transients and their effects on geospace and the

heliosphere. These questions are also pertinent, within the framework of the ‘solar-stellar

connection,’ to our understanding of other stellar systems that exhibit transient behavior such as

flaring (e.g., Getman et al. 2008).

Solar Orbiter will provide a critical step forward in understanding the origin of solar transient

phenomena and their impact on the heliosphere. Located close to the solar sources of transients,

Solar Orbiter will be able both to determine the inputs to the heliosphere and to measure directly the

heliospheric consequences of eruptive events at distances close enough to sample the fields and

plasmas in their pristine state, prior to significant processing during their propagation to 1 AU. Solar

Orbiter will thus be a key augmentation to the chain of solar-terrestrial observatories in Earth orbit

and at the libration points, providing a critical perspective from its orbit close to the Sun and out of

the ecliptic.

In the following sections, we discuss in more detail three interrelated questions which flow down

from this top-level question:

How do CMEs evolve through the corona and inner heliosphere? How do CMEs contribute to solar

magnetic flux and helicity balance? How and where do shocks form in the corona and inner

heliosphere?

3.2.2.1 2.1 How do CMEs evolve through the corona and inner heliosphere?


Following earlier observations by space-based white-light coronagraphs, considerable progress in

understanding CMEs has been achieved using data from the SOHO mission, which provides

continuous coverage of the Sun and combines coronagraphs with an EUV imager and off-limb

spectrometer. Other spacecraft, such as ACE, WIND, Ulysses and STEREO, which carried

comprehensive in-situ instrumentation, have contributed significantly to our understanding of the

interplanetary manifestation of these events. With a full solar cycle of CME observations, the basic

features of CMEs are now understood. CMEs appear to originate from highly-sheared magnetic

field regions on the Sun known as lament channels, which support colder plasma condensations



known as prominences. Eruptions develop in the low corona within 10-15 minutes, while the

associated shocks cross the solar disk within 1 hour. CMEs reach speeds of up to 3000 km/s and

carry energies (kinetic, thermal and magnetic) of ~1025 J (= 1032 ergs). They can also accelerate

rapidly during the very early stages of their formation, with the CME velocity being closely tied, in

time, to the associated flare’s soft X-ray light profile (Zhang and Dere 2006). High-resolution

SOHO and STEREO coronagraph images have provided evidence for a magnetic flux rope structure

in some CMEs as well as for post-CME current sheets. Both features are predicted by CME

initiation models (e.g., Lin & Forbes 2000; Lynch et al. 2004).

STEREO observations have made it possible to chart in three dimensions the trajectories of CMEs

in the corona and heliosphere, thereby improving our understanding of CME evolution and

propagation. STEREO data have supported detailed comparison both of in-situ measurements with

remote-sensing observations and of MHD heliospheric simulations with observations. The

combination of high-cadence coronagraphic and EUV imaging simplifies the separation of the CME

proper from its effects in the surrounding corona (Patsourakos and Vourlidas 2009) and allows a

more accurate determination of its dynamics.

Despite the advances in our understanding enabled by SOHO and STEREO, very basic questions

remain unanswered. These concern the source and initiation of eruptions, their early evolution, and

the heliospheric propagation of CMEs. All current CME models predict that the topology of ICMEs

is that of a twisted flux rope as a result of the flare reconnection that occurs behind the ejection.

Observations at 1 AU, however, find that less than half of all ICMEs, even those associated with

strong flares, have a flux rope structure (Richardson and Cane 2004). Many ICMEs at 1 AU appear

to have a complex magnetic structure with no clearly-defined topology. Moreover, for ICMEs that

do contain flux ropes, the orientation is often significantly different from that expected on the basis

of the orientation of the magnetic fields in the prospective source region. CMEs are believed to

originate from prominence eruptions, yet in ICMEs observed at 1 AU prominence plasma is very

rarely detected. These major disconnects between theoretical models (of prominence eruption and

CME propagation) and observations (remote and in situ) need to be resolved if any understanding of

the CME process is to be achieved.


To advance our understanding of the structure of ICMEs and its relation to CMEs at the Sun beyond

what has been achieved with SOHO and STEREO requires a combination of remote-sensing and in-

situ measurements made at close perihelion and in near-corotation with the Sun. Through combined

observations with its magnetograph, imaging spectrograph, and soft X-ray imager, Solar Orbiter

will provide the data required to establish the properties of CMEs at the Sun and to determine how

coronal magnetic energy is released into CME kinetic energy, flare-associated thermal/non-thermal

particle acceleration, and heating. Observations with the imaging spectrograph will be used to

determine the composition of CMEs in the low corona and to establish how they expand and rotate

and will also provide vital clues to the energy partition within a CME once it is released. Solar

Orbiter will make comprehensive in-situ measurements of the fields and plasmas (particularly

composition) of ICMEs following their release and, critically, prior to their processing during

propagation in the heliosphere. These measurements will allow the properties of an ICME to be

related to those of the CME at the Sun and to the conditions in the CME source region as observed

by Solar Orbiter’s remote-sensing instruments and will make it possible to examine the evolution of



CMEs in the inner heliosphere. Solar Orbiter’s combination of remote-sensing and in-situ

observations will also establish unambiguously the magnetic connectivity of the ICME and reveal

how the magnetic energy within flux ropes is dissipated to heat and accelerate the associated

particles. Solar Orbiter data will also reveal how the structure of the magnetic field at the front of a

CME evolves in the inner heliosphere – a critical link in understanding, and eventually predicting,

how transient events on the Sun may determine the geoeffective potential of the event.

To fully understand the physical system surrounding CME ejection, the temporal evolution of active

regions and CME-related shocks and current sheets must be tracked from their formation in the

corona to their expulsion in the solar wind. During the mission phases when the spacecraft is in

near-corotation with the Sun, Solar Orbiter will continuously observe individual active regions, free

from projection complications, over longer periods than are possible from Earth orbit. Solar Orbiter

will thus be able to monitor the development of sheared magnetic fields and neutral lines and to

trace the flux of magnetic energy into the corona. Observations from the vantage point of near-

corotation will make it possible to follow the evolution of the current sheet behind a CME with

unprecedented detail and to clarify the varying distribution of energy in different forms (heating,

particle acceleration, kinetic).

2.1.1 What are the global structure, initiation, and evolution of CMEs?

2.1.1.1 CME initiation


CME initiation has been a core space physics problem for the last three decades. The two current

paradigms are distinguished primarily by the topology of the pre-eruption magnetic field: twisted

flux rope (e.g., Roussev et al., 2004) or sheared arcade (Antiochos et al., 1994). Irrespective of the

pre-eruption topology, all models predict that as a result of the flare reconnection occurring below

the ejection, CMEs in the heliosphere must have a twisted flux rope topology, as commonly

observed (Gosling et al., 1995). If the pre-eruption topology is that of a twisted flux rope, then the

innermost part of its structure should exhibit relatively undisturbed filament plasma parameters.

However, if the twist forms only as a result of flare reconnection, then the whole twisted structure in

the heliosphere should exhibit the properties of flare-reconnection-heated plasma, hot beamed

electrons, high charge states of Fe, as well as compositional anomalies of heavy ions including He.

By measuring the electron and ion properties of a CME along with its magnetic structure, we

determine the pre-eruption topology and the initiation mechanism. Solar Orbiter will provide the

opportunity to perform these measurements near the Sun, minimizing propagation effects such as

internal reconnection, which homogenizes the CME structure (e.g., Lynch et al., 2005).

2.1.1.2 CME structure


In order to better understand the structure of CMEs, we have to explore the following:



• Three-part structure: bright leading shell, dark cavity, bright core (Illing and Hundhausen,

1986; Plunkett et al., 2000).

• Presence of large-scale helical magnetic flux rope in the CME cavity (Vourlidas et al.,

2013).

• Counter-streaming electrons observations as signatures of field lines tied at both ends to the

Sun (Crooker and Owens, 2012).

• Electron heat flux dropouts as signatures of magnetic flux disconnection from the Sun

(Crooker and Owens, 2012).

• Longitudinal distribution of the emitted plasma and detection of geoeffective events.

2.1.1.3 CME evolution


In order to better understand the structure of CMEs, we have to explore the following:

• Self-similar expansion (Cremades and Bothmer, 2004; Gibson and Low, 1998).

• Cylindrical geometry with the axis of symmetry corresponding to the long axis of a large-

scale helical magnetic flux rope that originates in the CME source region (Thernisien et al.,

2006).

• Interaction with the background velocity and density structures (Odstrcil et al., 2005; Riley

et al., 2003; Colaninno et al., 2014).

• Aerodynamic drag force and equalization of ICME and solar wind speeds (Gosling and

Riley, 1996; Cargill, 2004; Vrsnak et al., 2012; Subramanian et al., 2012).

• Relationship between the three-part structure of a CME and the ICME counterparts

(magnetic cloud with a flux rope or ‘complex ejecta’ with disordered magnetic fields)

(Vourlidas et al., 2013).

• Bright front evolution to become the sheath of compressed solar wind.

• Dark cavity corresponds to the flux rope.

• Presence of an in situ ‘plug’ of cold, dense plasma trailing the flux rope, interpreted as

remnant material from the erupting filament (bright core).

• What is the fate of the erupting filament as the CME propagates?

o Only a small fraction escapes with the CME?

o Filament material flows back along magnetic field lines or falls back due to

Rayleigh-Taylor instability (Innes et al., 2012; Carlyle et al.,2014; van Driel-

Gesztelyi et al., 2014).

o Filament is present but has lost its expected low-charge state signature because of

heating (Skoug et al., 1999; Rakowski et al., 2007)?

• Relationship between the three-part structure of a CME and shocks in the low corona and in

situ.

• Interaction between CMEs and its effects on SEPs (Gopalswamy et al., 2004; Richardson et

al., 2003; Rodriguez-Pacheco et al., 2003).

• Dynamics in transient coronal holes and recovery phase of the eruption.

• Improve CME arrival time estimates and predictions of geomagnetic activity.



3.2.2.2 2.2 How do CMEs contribute to solar magnetic flux and helicity balance?


Magnetic flux is transported into the heliosphere both by the solar wind, in the form of open flux

carried mostly by the fast wind from polar coronal holes, and by coronal mass ejections, which drag

closed flux with them as they propagate into the heliosphere. At some point, the closed flux

introduced by CMEs must be opened to avoid an unsustainable buildup of magnetic flux in the

heliosphere. Measurements of the magnetic flux content of the heliosphere from near the Earth

show that the total amount of magnetic flux in the solar system changes over the solar cycle.

Longer-term variations are also known to occur. Proxies such as geomagnetic activity and cosmic

ray fluxes provide evidence that the average IMF strength has increased substantially in the last 100

years, perhaps by as much as a factor of 2. Surprisingly, however, during the recent solar minimum,

the IMF strength is lower than at any time since the beginning of the space age.

The relative contribution of the solar wind and CMEs to the heliospheric magnetic flux budget is an

unresolved question, as is the process by which the flux added by the CMEs is removed. Models to

explain the solar cycle variation assume a background level of open flux, to which CMEs add extra

flux during solar maximum, increasing the intensity of the IMF. The exceptionally low intensity of

the IMF during the last minimum has been attributed to the low rate of CME occurrence [Owens et

al. 2008]. Alternatively, there may simply be no ‘background’ open flux level.

There is evidence that the flux introduced into the heliosphere by CMEs may be removed by

magnetic reconnection within the trailing edges of CMEs, which disconnects the CME from the Sun

or by interchange reconnection closer to the solar surface [e.g., Owens and Crooker 2006]. Recent

observations show that the reconnection process occurs quite often in the solar wind, even when the

magnetic field is not under compression. However, the rate and/or locations at which reconnection

generally removes open flux are not at present known.

Together with magnetic flux, the solar wind and CMEs carry magnetic helicity away from the sun.

Helicity is a fundamental property of magnetic fields in natural plasmas, where it plays a special

role because it is conserved not only by the ideal dynamics but also during the relaxation which

follows instabilities and dissipation. Helicity is injected into the corona when sunspots and active

regions emerge, via the twisting and braiding of magnetic flux. During the coronal heating process,

the overall helicity is conserved and tends to accumulate at the largest possible scales. It is natural to

assume that critical helicity thresholds may be involved in the triggering of CMEs, but how solar

eruptions depend on the relative amounts of energy and helicity injection during active region

emergence and evolution is unknown. Yet this understanding could be a crucial element in the

prediction of large solar events.


Fundamental to the question of the contribution of CMEs to the heliospheric flux budget is the flux

content of individual events. Encountering CMEs close to the Sun before interplanetary dynamics



affect their structure, Solar Orbiter will measure their magnetic flux content directly; comparisons

with remote-sensing measurements of their source regions will clarify the relation between CME

flux and the eruption process. As Solar Orbiter moves through the inner heliosphere, it will

encounter CMEs at different solar distances, making it possible to quantify the effect of

interplanetary dynamics on their apparent flux content.

The flux carried outwards by CMEs must eventually disconnect completely from the Sun, or

interchange reconnects with existing open field lines. Solar Orbiter will diagnose the magnetic

connectivity of the solar wind and CME plasma using suprathermal electron and energetic particle

measurements. These particles, which stream rapidly along the magnetic field from the Sun, indicate

whether a magnetic flux tube is connected to the Sun at one end, at both ends, or not at all. These

particles disappear when the field is completely disconnected, or may reverse their flow direction as

a result of interchange reconnection. However, scattering and reflection due to curved, tangled, or

compressed magnetic field lines act to smear out these signatures with increasing solar distance,

leading to ambiguity in connectivity measurements. Solar Orbiter, by traveling close to the Sun

before this scattering is significant, will determine the original level of magnetic connectivity;

covering a wide range of distances in the inner heliosphere, the spacecraft will measure how the

connectivity changes as field lines are carried away from the Sun.

Solar Orbiter will also directly sample reconnection regions in the solar wind as they pass the

spacecraft, determining their occurrence rates in the inner heliosphere as a function of distance and

testing theories of CME disconnection by searching for reconnection signatures in the tails of

CMEs.

The contribution to the heliospheric magnetic flux of small scale plasmoids, ejected from the tops of

streamers following reconnection events, is unclear. Solar Orbiter, slowly moving above the solar

surface during perihelion passes, will determine the magnetic structure, connectivity, and plasma

properties including the composition of these ejecta, using spectroscopic imaging observations to

unambiguously link them to their source regions.

To assess the role of CMEs in maintaining the solar magnetic helicity balance, Solar Orbiter will

compare the helicity content of active regions as determined from remote sensing of the

photospheric magnetic field with that of magnetic clouds measured in situ. Such a comparison

requires both extended remote-sensing observations of the same active region over the region’s

lifetime and in-situ measurements of magnetic clouds from a vantage point as close to the solar

source as possible. Around its perihelia, Solar Orbiter will ‘dwell’ over particular active regions and

observe the emergent flux for a longer interval (more than 22 days) than is possible from 1 AU,

where perspective effects complicate extended observations. The resulting data will be used to

calculate the helicity content of an active region, track its temporal variation, and determine the

change in helicity before and after the launch of any CMEs. Should a magnetic cloud result from an

eruptive event in the active region over which Solar Orbiter is dwelling, the relatively small

heliocentric distances between the solar source and the spacecraft will make it highly probable that

Solar Orbiter will directly encounter the magnetic cloud soon after its release. Determination of the

cloud’s properties and connectivity through Solar Orbiter’s in-situ particle-and-fields measurements

will enable the first-ever comparison of a magnetic cloud in a relatively unevolved state with the

properties of the solar source, an impossibility with measurements made at 1 AU. The comparison

of the helicity change in the source region with the value measured in the magnetic cloud will

provide insight into the role of CMEs in the helicity balance of the Sun.



2.2.1 How do CMEs contribute to the global evolution of magnetic flux in the heliosphere?


• Temporal addition of closed flux by CMEs, which opens via interchange reconnection,

conserves open magnetic flux while accounting for the rise in the observed heliospheric field

strength near solar maximum (Owens and Crooker, 2006).

• CME induced interchange reconnection may also drive the reversal of the heliospheric open

field (Owens et al., 2006).

• For this we need to understand:

• How much axial magnetic flux is carried out by CMEs?

• The opening time of closed CME flux through interchange reconnection.

• The latitudinal separations of CME foot-points.

• How does the in situ ICME magnetic flux compare to photospheric magnetic flux?

2.2.2 What is the role of ICMEs in the Sun’s magnetic cycle?


• Measure the helicity of ICMEs leaving the Sun and link them directly to changes in the solar

field thereby quantifying the effect of ICMEs on the solar cycle (Owens et al., 2007).

• Study the evolution of ICMEs in the inner heliosphere to determine the effects of local

dynamics in changing their characteristics and apparent helicity, relating them directly to

observed solar events (Zurbuchen and Richardson, 2006).

3.2.2.3 2.3 How and where do shocks form in the corona?


The rapid expulsion of material during CMEs can drive shock waves in the corona and heliosphere.

Shocks in the lower corona can also be driven by flares, and in the case of CME/eruptive flare

events, it may be difficult to unambiguously identify the driver (Vršnak and Cliver 2008). CME-

driven shocks are of particular interest because of the central role they play in accelerating coronal

and solar wind particles to very high energies in SEP events.

Shocks form when the speed of the driver is super-Alfvénic. The formation and evolution of shocks

in the corona and the inner heliosphere thus depend (1) on the speed of the driving CME and (2) on

the Alfvén speed of the ambient plasma and its spatial and temporal variations. According to one

model of the radial distribution of the Alfvén speed in the corona near active regions, for example,

shocks can form essentially in two locations, in the middle corona (1.2-3 RSun), where there is an

Alfvén speed minimum, and distances beyond an Alfvén speed maximum at ~4 RSun (Mann et al.



2003). A recent study of CMEs with and without type II radio bursts (indicative of shock formation)

has shown that some of the fast and wide CMEs observed produced no shock or only a weak shock

because they propagated through tenuous regions in the corona where the Alfvén velocity exceeded

that of the CME (Gopalswamy et al. 2008). CME shock formation/evolution can also be affected by

the interaction between an older, slower-moving CME and a faster CME that overtakes it.

Depending on the Alfvén speed in the former, the interaction may result in the strengthening or

weakening of an existing shock driven by the overtaking CME or, if there is no existing shock, the

formation of one (Gopalswamy 2001; 2002).

Studies of LASCO images obtained during the rising phase of solar cycle 23 have demonstrated the

feasibility of detecting CME-driven shocks from a few to ~20 RSun and of measuring their density

compression ratio and propagation direction (Vourlidas et al. 2003; Ontiveros and Vourlidas 2009).

This development has opened the way for the investigation of shock formation and evolution in the

lower corona and heliosphere through Solar Orbiter’s combination of remote-sensing observations

and in-situ measurements.


Understanding shock generation and evolution in the inner heliosphere requires knowledge of the

spatial distribution and temporal variation of plasma parameters (density, temperature, and magnetic

field) throughout the corona. Solar Orbiter’s remote-sensing measurements – in particular electron

density maps derived from the polarized visible-light images and maps of the density and outflow

velocity of coronal hydrogen and helium – will provide much improved basic plasma models of the

corona, so that the Alfvén speed and magnetic field direction can be reconstructed over the distance

range from the Sun to the spacecraft. Remote sensing will also provide observations of shock

drivers, such as flares (location, intensity, thermal/non-thermal electron populations, time-profiles),

and manifestations of CMEs (waves, dimmings, etc.) in the low corona with a spatial resolution of a

few hundred kilometers and cadence of a few seconds. It will measure the acceleration profile of the

latter and then track the CMEs through the crucial heights for shock formation (2-10 RSun) and

provide speed, acceleration, and shock compression ratio measurements.

Type II bursts, detected by Solar Orbiter, will indicate shock-accelerated electron beams produced

by the passage of a CME and thus provide warning of an approaching shock to the in-situ

instruments. These in-situ plasma and magnetic field measurements will fully characterize the

upstream and downstream plasma and magnetic field properties and quantify their microphysical

properties, such as turbulence levels and transient electric fields (while also directly measuring any

SEPs). Spacecraft potential measurements also allow for rapid determinations of the plasma density,

and of electric and magnetic field fluctuations, on microphysical scales, comparable to the Doppler-

shifted ion scales, which are characteristic of the spatial scales of shocks. The evolution of such

parameters will provide insight into the processes dissipating shock fronts throughout the range of

magnetic/velocity/density and pressure parameter space. Because of Solar Orbiter’s close proximity

to the Sun, the measurements of the solar wind plasma, electric field, and magnetic field will be

unspoiled by the dynamical wind interaction pressure effects due to solar rotation and will provide

the first reliable data on the magnetosonic speed, the spatial variation of the plasma pressure and

magnetic field in the inner heliosphere. MHD modeling studies have shown that interactions among

recurring CMEs and their shocks occur typically in the distance range around 0.2-0.5 AU (Lugaz et

al. 2005). Solar Orbiter will spend significant time in the regions of recurring CME interactions and



so will be able to investigate the effects of such interactions on the evolution of CME-driven

shocks.

2.3.1 Coronal shocks


• Identify coronal shocks and characterize their spatial distribution and outward propagation

velocity.

• Study interaction with coronal plasma.

• Characterize the longitudinal distribution of coronal shocks during high latitude orbits.

2.3.2 What are the properties and distribution of heliospheric shocks?

2.3.2.1 Understand coronal conditions under which the shocks form and determine the

interplanetary conditions where they evolve

2.3.2.2 Identify interplanetary shocks and characterize their spatial and temporal evolution

• Identify interplanetary shocks and characterize their spatial and temporal evolution.

• Characterize shock structure and determine whether they are locally quasi-perpendicular or

quasi-parallel, and study kinetic properties of shock-related waves and turbulence that

control ion scattering mean free paths near the shock.

2.3.2.3 Study heating and dissipation mechanisms at shocks with radial distance

2.3.2.4 Identify mechanisms that heat the thermal solar wind particle populations near shocks and

determine their energy partition

2.3.3 Resolve the interplanetary shock field and plasma structure down to the spatial and temporal

scales comparable and smaller than the typical ion scales.

2.3.4 Shock-surfing acceleration mechanism


• Measure the relative strength of the convection vs cross-shock electric field as a function of

heliocentric distance.

• Fine-scale structure of shocks with fast density measurements (using satellite potential

measurements).

2.3.5 Understand the radio emissions from the ICME driven shocks




• Type II and III radio bursts.

• Langmuir waves and electromagnetic mode conversion.

• Characterize the energy balance between electron beams, Langmuir waves and

electromagnetic radio waves at several heliocentric distances.

2.3.6 Identify shock accelerated particles

• Infer the rates of particle acceleration and the injection energies as a function of distance and

along different parts of the shock.



3.2.3 Objective 3: How do solar eruptions produce energetic particle radiation that fills the

heliosphere?

Astrophysical sites throughout the solar system and galaxy have the universal ability to

accelerate ions and electrons to high speeds, forming energetic particle radiation. Detected remotely

from radio and light emission around supernovae remnants, the Sun, and planets, or directly from

particles that reach our detectors, this radiation arises from the explosive release of stored energy

that can cause magnetic fields to rearrange, or can launch shock waves which accelerate particles by

repeatedly imparting many small boosts to their speed. The nearly universal occurrence of energetic

particle radiation, along with the effects it can have on planetary environments, evolution of life

forms, and space systems has fostered a broad interest in this phenomenon that has long made it a

high priority area of investigation in space science. Since remote sites in the galaxy cannot be

studied directly, solar system sources of energetic particles give the best opportunity for studying all

aspects of this complex problem.

The Sun is the most powerful particle accelerator in the solar system, routinely producing

energetic particle radiation at speeds close to the speed of light, sufficiently energetic to be detected

at ground level on Earth even under the protection of our magnetic field and atmosphere. SEP

events can severely affect space hardware, disrupt radio communications, and cause re-routing of

commercial air traffic away from polar regions. In addition to large events, which occur roughly

monthly during periods of high sunspot count, more numerous, smaller solar events can occur by

thousands each year, providing multiple opportunities to understand the physical processes

involved. In the following sections, we have divided in more detail three interrelated questions that

flow down from this top-level question: How and where are energetic particles accelerated at the

Sun? How are energetic particles released from their sources and distributed in space and time?

What are the seed populations for energetic particles?

3.2.3.1 3.1 How and where are energetic particles accelerated at the Sun?

One of the two major physical mechanisms for energizing particles involves particles

interacting with moving or turbulent magnetic fields, gaining small amounts of energy at each

step and eventually reaching high energies. Called Fermi or stochastic acceleration, this mechanism

is believed to operate in shock waves and in turbulent regions such as those associated with

reconnecting magnetic fields or in heated coronal loops. The second major physical mechanism is a

magnetic field whose strength or configuration changes in time, producing an electric field which

can directly accelerate particles in a single step. At the Sun, such changes occur when large

magnetic loop structures reconnect or are explosively rearranged due to the stress from the motion

of their footpoints on the solar surface (e.g., Aschwanden 2006; Giacalone and Kota 2006).

Multiple processes may take place in SEP events, and while it is not possible to cleanly

separate them, they can be split into two broad classes, the first being events associated with shock

waves. As a CME moves into space, it drives a shock creating turbulence that accelerates SEPs from

a seed population of ions filling the interplanetary medium. Mixed into this may be particles from



an associated solar flare. CMEs often accelerate particles for hours as they move away from the

Sun, and in some cases, are still accelerating particles when they pass Earth orbit in a day or two.

Since CMEs can be huge, it is easy to see how they can fill a large portion of the heliosphere with

SEPs. However, the correlation of the observed radiation intensities with CME properties is poor,

indicating that additional aspects of the mechanism such as seed populations or shock geometry

must play important roles that are not yet well understood (Gopalswamy 2006; Desai et al. 2006;

Mewaldt 2006).

The second class of events is associated with plasma and magnetic field processes in loops

and active regions that accelerate particles. Reconnecting magnetic loops, and emerging magnetic

flux regions provide sites for stochastic energetic particle acceleration or acceleration by electric

fields. Because these regions are relatively small, the acceleration process is quick: on the order of

seconds or minutes, but the resulting event is small and often difficult to observe. Since the

energized particles are in the relatively high-density regions of the corona, they collide with coronal

plasma, producing ultraviolet (UV) and X-ray signatures that make it possible to locate their

acceleration sites and probe the local plasma density. Most of these particles remain trapped in their

parent loops, traveling down the legs to the solar surface where they lose their energy to the ambient

material, producing X- and gamma-rays. A few escape on magnetic field lines leading to

interplanetary space, traceable by their (‘type III’) radio signatures, electrons, and highly

fractionated ion abundances where the rare 3He can be enhanced by 1000-10,000 times more than in

solar material.

The energetic particles from these events reach our detectors at Earth orbit after spiraling

around the IMF, which is an Archimedes spiral on average. But since the IMF meanders, and has

many kinks, the length of the particle’s path has a good deal of uncertainty, and the particles

themselves scatter and mix, smearing and blurring signatures of the acceleration at the Sun.

Although we can enumerate candidate mechanisms for producing SEPs, a critical question is: what

actually happens in nature? Which processes dominate? How can shocks form fast enough to

accelerate ions and electrons to relativistic energies in a matter of minutes?

3.1.0 Explore in depth the SEP properties

Our current knowledge of the properties of the two different kinds of SEP events, gradual and

impulsive can be summarized as follows:

Gradual SEP events:

▪ electrons up to tens of MeV, ions at GeV

▪ once per month near solar maximum

▪ dominated by protons (small e/p fluence ratios)

▪ variable composition and charge states

▪ extend over > 100° in solar longitude

▪ generally associated with large solar flares & fast CMEs

▪ believed to be due to acceleration by CME-driven shocks and not by flares.

Impulsive SEP events:



• Low energy electrons: 1-100 keV

• Ions: 0.01-1 MeV/nucleon

• >104 events per year near solar maximum

• large e/p ratios

• enhanced abundances of 3He and of heavy ions

• high charge states

• enhanced alpha/proton ratio

• extend over 30° in solar longitude

• associated with type III bursts

• many clearly come from flares/microflares, but often not even a microflare or soft X-ray

burst is observed.

3.1.1 CME and shock associated SEP sources

Moving from the lower corona to the interplanetary medium, shocks evolve rapidly since the

sound speed drops as plasma density and magnetic field strength decline as ~1/r^2. Solar Orbiter’s

coronagraphs will remotely identify shock front location, speed, and compression ratios through

this critical region within ~10 RSun. Combining this information with local electron densities as

well as coronal ion velocities given by Solar Orbiter radio and light polarization observations will

provide critical constraints on shock evolution models in regions too close to the Sun for direct

sampling.

In the regions explored by Solar Orbiter close to the Sun, the IMF is almost radial with much

less variation (uncertainty) in length than is the case at 1 AU, so the knowledge of the actual path

length improves by a factor of 3-5 as the length shortens. Having observed the CMEs and their radio

signatures in the corona and the X-ray signatures of the energetic particles near the Sun, Solar

Orbiter will then determine the subsequent arrival time of the particles in situ that can be accurately

compared to CME position. As the shock then rolls past the spacecraft, Solar Orbiter will measure

the shock speed and strength as well as the associated plasma turbulence, electric, and magnetic

field fluctuations. This will give a complete description of the acceleration parameters in the inner

heliosphere where much of the particle acceleration takes place. Indirect evidence from 1 AU

indicates that shock acceleration properties depend on the longitude of the shock compared to

the observer; close to the Sun, Solar Orbiter can cleanly test this property since the IMF is nearly

radial, the CME lift-off site is known, and the accelerated particles will have little chance to mix. In

the high-latitude phase of the mission, Solar Orbiter will be able to look down on the longitudinal

extent of CMEs in visible, UV, and hard X-rays, allowing first direct observations of the

longitudinal size of the acceleration region. This will make it possible to test currently

unconstrained acceleration and transport models by using measured CME size, speed, and shape to

specify the accelerating shock.

3.1.1.1 Where and when are shocks most efficient in accelerating particles?



Type II radio bursts show that CME-driven shocks form in the corona, but we do not know

when they start accelerating ions. We do know from energetic storm particle (ESP)/shock spike

events that shocks accelerate particles at 1 AU, but rarely to >~ a few 10s of MeV. As the shock

propagates and evolves, when is it most potent in accelerating particles? How is the acceleration

related to the changing magnetic geometry? Comparisons of the kinetic energy of fast CMEs

measured by SOHO/LASCO with ACE and GOES measurements of accelerated particle energy

spectra indicate that as much as 10% of the CME kinetic energy often goes into accelerated

particles. That is about the efficiency needed for the acceleration of galactic cosmic rays by

supernovae shocks.

Correlating the SEP/CME efficiency estimates with interplanetary conditions near the Sun

will probe why and how the CME acceleration efficiency is so variable. With the wide range of

other multi-point imaging and in situ data available, it will be possible to significantly improve on

estimates of the CME energy, and of the global energy budget for solar eruptions, including plasma,

magnetic field, energetic particle, and photon contributions. (e.g., Emslie et al. 2004).

There is now a well-developed theory of SEP shock acceleration and the capability to

perform sophisticated simulations of these events. The models, however, depend on assumptions

about conditions in the inner heliosphere, and many predictions of the models cannot be verified

using data from 1 AU. By going close to the Sun, within ~1-2 λ, SolO provides a unique opportunity

to probe shock acceleration of SEPs to high energies with powerful new in situ and imaging

diagnostic measurements. Only by measuring the accelerated particles and the plasma and magnetic

field properties close to the acceleration source can we test and improve current models of the SEP

acceleration and transport.

We, therefore, need to explore the following questions:

• As the shock propagates and evolves, when is it most potent in accelerating particles?

o When/where does the shock start accelerating ions?

• How is the acceleration related to the changing magnetic geometry?

• Improve SEP kinetic energy estimates (variable CME acceleration efficiency) (Mewaldt et

al., 2005)

o As much as 10% of the CME kinetic energy often goes into accelerated particles.

o Show why the efficiency is so variable (by better constraining CME energies?).

• What is the relation between shock acceleration, turbulence properties, and anomalous

diffusion mechanisms? (Perri and Zimbardo, 2012)

3.1.1.2 Why are gradual SEP events so variable?

One of the most intriguing questions about gradual events is their great variability in

intensity: while the CME speed varies by a factor of 4, the SEP intensities of >20 MeV protons vary

by 3 to 4 orders of magnitude for a given CME speed (Kahler and Vourlidas 2005). In addition,

there are wide variations in the composition of gradual SEP events, including Fe/O ratios that vary

by a factor of ~100 (Mewaldt et al. 2006).

Observations from solar cycle 23 suggest that CME-driven shocks accelerate mainly

suprathermal ions rather than the bulk solar wind (Mason et al. 1999). The intensity and

composition of suprathermal ions with 10 to 100 keV/nuc is known to be much more variable than



the solar wind. One reason for this variability is that suprathermal particles are believed to arise

from a variety of sources, including previous gradual and impulsive SEP events, ubiquitous micro-

and nano-flaring at the Sun, and interplanetary acceleration of solar wind by stochastic processes

(Fisk and Gloeckler 2006). In a steady state solution, particles are assumed to be continually

injected at the shock and a power law spectrum is achieved when there is enough time. However, at

a propagating CME-driven shock, the amount of time available for acceleration is limited. Thus, for

the acceleration to successfully operate, it is crucial that there are enough seed particles.

3.1.1.2.1 Intensity variability

• Study and explain the SEP intensity variability (Kahler and Vourlidas, 2005): while the

CME speed varies by a factor of 4, the SEP intensities of >20 MeV protons vary by 3 to 4

orders of magnitude for a given CME speed.

3.1.1.2.2 Composition variations

The intensity and composition of suprathermal ions with 10 to 100 keV/nuc is known to be

much more variable than the solar wind. One reason for this variability is that suprathermal particles

are believed to arise from a variety of sources, including previous gradual and impulsive SEP

events, ubiquitous micro- and nano-flaring at the Sun, and interplanetary acceleration of solar wind

by stochastic processes (Fisk and Gloeckler 2006).

In a steady state solution, particles are assumed to be continually injected at the shock and a

power law spectrum is achieved when there is enough time. However, at a propagating CME-driven

shock, the amount of time available for acceleration is limited. Thus, for the acceleration to

successfully operate, it is crucial that there are enough seed particles.

We need observations of the suprathermals and SEPs before and during CMEs, to know

what material is actually accelerated. Simultaneous observations of the near Sun seed particles with

SPP would be very helpful.

3.1.1.2.3 Warped shock fronts

3.1.1.2.4 Turbulence and inhomogeneities

3.1.1.3 How are energetic particles accelerated continuously in the corona and solar wind?

Observations of quiet-time suprathermal ions and electrons (the "superhalo") at 1 AU and

beyond find that the suprathermal tails of the solar wind are always present, suggesting that

continuous acceleration is taking place in the solar wind and/or in the corona. Fisk and Gloeckler

(2006) have suggested that the ubiquitous suprathermal ion tails can result from thermal particles

interacting with compressional turbulence embedded in the solar wind. This mechanism predicts a

power-law steady state spectrum with an index of -1.5, consistent with a number of ion observations

at 1 AU and in the outer heliosphere. However, the mechanism of the acceleration takes time, so it

is likely that the suprathermal spectrum will not be fully developed and more variable inside 1 AU.



3.1.1.4 How can SEPs be accelerated to high energies so rapidly?

The January 20, 2005 event provided a wake-up call to space weather forecasters when the

>100 MeV proton intensities reached their highest levels in more than 15 years. This event was also

the 2nd largest ground-level event of all time, reaching peak intensity only 18 minutes after the

onset of the associated X-ray flare. If this event was accelerated by a CME-driven shock, the shock

must have formed very low in the corona and accelerated particles to GeV energies within ~10

minutes or less (Mewaldt et al. 2005a; Saiz et al. 2005), presenting a significant challenge to shock

acceleration models. On the other hand, Simnett (2006) and others have suggested that the January

20 SEPs were accelerated by the X7 flare.

Shock geometry may be a key parameter for understanding how particles can be rapidly

accelerated to very high energies. According to simulations by Giacalone (2005), particles can be

accelerated much more rapidly at quasi- perpendicular shocks than they can at quasi-parallel shocks.

Although it has also been claimed that quasi-perpendicular shocks have a higher injection threshold

than quasi-parallel shocks and therefore favor the acceleration of pre-existing suprathermal ions

over thermal seed particles (Tylka et al. 2005), this is a subject of controversy (Giacalone 2005).

These controversies are best tested close to the Sun, where particles are accelerated to higher

energies than at 1 AU.

We, therefore, have to test whether:

• there is a correlation between shock geometry and the composition and energy spectra of

particles when the shock is still close to the Sun (Giacalone 2005; Mewaldt et al., 2005; Saiz

et al. 2005),

• anomalous diffusion mechanisms can give shorter acceleration times (Zimbardo and Perri,

2013).

3.1.1.5 Do proton-amplified Alfvén waves play a role in accelerating particles at shocks?

At quasi-parallel shocks, the diffusive shock acceleration mechanism operates when

particles repeatedly cross the shock front and gain energy by scattering off magnetic irregularities

carried by the plasma flow. One of the essential elements of this theory is that protons streaming

upstream from the shock can amplify anti-sunward propagating Alfvén waves in the solar wind. The

amplified waves can then resonantly scatter subsequent ions escaping upstream, trapping them near

the shock and increasing the efficiency of the shock acceleration process (Lee 1983). Although the

existence of these waves is a cornerstone of the theory (also invoked for galactic cosmic ray

acceleration by supernovae shocks), the proton beams and the amplified waves that are required for

diffusive shock acceleration of SEPs to high energies (>~10-100MeV) at a CME shock have never

been observed.



The waves responsible for accelerating MeV ions at shocks are seldom observed at 1 AU,

where they are difficult to observe against the background turbulence (e.g., Bamert et al. 2004), but

in the inner heliosphere, wave growth is expected to proceed more rapidly, and the wave intensities

are predicted to be significantly greater (Ng, et al, 2003) and should be easily detectable at 0.35 AU.

Solar Orbiter observations of the amplified waves by the magnetometer (MAG) and RPW search

coil (for electrons), in association with the proton (and electron) beams and simultaneous

accelerated particle spectra (by EPD), would be a crucial test of this fundamental theory.

3.1.1.6 What causes SEPs’ spectral breaks?

The energy spectra of most gradual SEP events start out as power laws at low energy with a

gradual spectral steepening in the energy range from ~3 to 30 MeV, indicating that the process has

become less efficient at high energies. The “spectral breaks” of heavier species occur at lower

energy/nucleon than for protons, suggesting a rigidity-dependent process (Tylka et al. 2001, Cohen

et al. 2005). It has been suggested that the breaks are organized by diffusion and that they occur at

the locations where there are sudden decreases in the wave intensity. Thus, the proton-amplified

Alfvén waves may also be responsible for determining the Q/A-dependence of spectral breaks in

large SEP events (Cohen et al. 2005, Mewaldt et al. 2005b, Li et al. 2005). Spectral breaks are also

seen in flare accelerated SEPs (Mason et al. 2002), so we require observations of the source as well.

3.1.1.7 Are there favorable environments for particle acceleration?

Statistical studies suggest that CMEs that erupt soon after a previous CME are more efficient

in accelerating particles than those erupting into a pristine environment (Gopalswamy et al. 2004).

Model calculations (Li and Zank, 2005) suggest that particles can be accelerated to energies ~30

times higher at a second shock following a previous shock. This more efficient acceleration may be

due to a stronger turbulence level and a larger population of the seed particles at the second shock,

but there are also other suggested explanations for the observations (see Gopalswamy et al. 2004).

3.1.2 SEPs associated with flares, coronal loops and reconnection regions

As Solar Orbiter approaches the Sun, the photon and particle signatures from small events

will increase by 1/r^2, making it possible to observe events 15-20 times smaller than ever before, in

effect opening a new window for SEP processes. We may detect for the first time energetic particle

populations from X-ray microflares, a candidate mechanism for coronal heating that cannot be

studied further away from the Sun due to background problems. For the small flares that produce X-

ray, electron, and 3He-enrichments we will observe with great accuracy events that at 1 AU are not

far above the level of detection: the timing of particle and radio signatures, the composition and

spectra, etc., providing strong new constraints on the process operating in these events. Particle

acceleration on coronal loops will have new insights since the 1/r^2 sensitivity advantage and



viewing geometry will make it possible to view the X-ray emission from the tops of loops in

numerous cases where the much stronger footpoint sources are occulted behind the solar limb.

These studies of faint coronal sources that are only rarely observable from 1 AU will give crucial

information about the location and plasma properties of suspected electron acceleration sites in the

high corona.

3.1.2.0 Impulsive SEP event sources

Small impulsive events dominated by ~1-100 keV (sometimes down to ~0.1 keV) solar

electrons are detected hundreds of times a year at ~1 AU near solar maximum (see Lin 1985). Over

the last solar cycle, the occurrence frequency rate for impulsive electron events observed at ~ 1AU

varies with the electron peak flux as a power law with average exponent -1.4, suggesting that many

more events are likely below present detection thresholds. Associated increases in ~0.01 to 1

MeV/nuc ions are often observed (may always be present, but with fluxes too low to be always

detected), with factors of 10-10000 enrichments of 3He (sometimes 3He/4He >1) along with factors

of ~10 enrichment of heavy ions up to Fe, together with high charge states. Curiously, analysis of

the nuclear lines and continuum for the large γ-ray flares that accelerate ions to >~100 MeV shows

that those flares often also have high e/p ratios, enhanced heavy ions, and enhanced He/p ratios

(Ramaty et al. 1993), even though impulsive SEP events at 1 AU are almost never associated with

such large flares.

3.1.2.1 Understand energy release and particle acceleration process.

Large solar flares rapidly accelerate ions up to many GeV and electrons up to 10s of MeV,

and they could be a significant contributor to the SEPs observed in space in gradual events. Imaging

in hard X-rays generated by accelerated electrons, and in the 2.223 MeV neutron-capture γ-ray line

produced by accelerated >~30 MeV ions, show that these emissions come from the footpoints of

newly formed loops (Hurford et al, 2006). Furthermore, the accelerated >~20 keV electrons and >~

few MeV ions often contain >~10-50% of the energy released in the flare (Lin & Hudson 1976; Lin

et al 2003), suggesting that the particle acceleration is intimately related to the magnetic

reconnection process that appears to release the flare energy. Neither the energy release or particle

acceleration processes are understood, but some accelerated particles could end up on magnetic field

lines connected to the heliosphere, as observed in many impulsive events.

Understand the energy release and particle acceleration process:

• Origin from the footprints of newly formed loops (Hurford et al., 2006).

• Acceleration relating to the magnetic reconnection process.

• Properties of flare energy release (in high spatial and temporal resolution).



3.1.2.2 Evaluate how significantly large flares contribute directly to gradual SEP events

For the magnetically well-connected 2 November 2003 and 20 January 2005 flares, the

spectra inferred for the energetic protons that produce the γ-ray lines observed by RHESSI are

essentially the same, within measurement uncertainties, as the spectra of SEP protons at 1AU. As

already mentioned, the extremely rapid rise of SEP fluxes at 1 AU after the 20 January 2005 flare

X-ray peak raises serious questions about CME-driven shock acceleration. Measurements of

SEPs close to the Sun will provide the timing and compositional information to evaluate how

significantly do large flare contribute directly to gradual SEP events.

3.1.2.3 Flare seed particles

For SEP events in general, the intensity of the photon emission appears poorly correlated

with the particle intensities observed in interplanetary space, so it is not clear whether many of these

particles escape. Particles accelerated by flares of all sizes down to microflares, however, may

provide the seed particles that are preferentially accelerated by interplanetary processes, thus

inseparably linking the two acceleration sites and processes. It has been proposed that SEP events

with large initial Fe/O ratios are a signature of “flare particles” followed by more typical Fe/O ratios

from subsequent interplanetary acceleration. Solar Orbiter observations can test these differing

models since its proximity to the flare site removes most of the uncertainty in magnetic connection,

and the timing differences between the flare acceleration and the interplanetary acceleration phase

would be much clearer than at 1 AU.

3.1.2.4 Explore the fact that only some of the hard X-ray peaks are related to escaping

electrons, while others are not (Benz et al., 2005).

3.1.2.5 X-ray prompt events

In X-ray literature "prompt event" is defined as an event where the estimated release time of

the electrons that are later seen at 1 AU agrees with the hard X-ray flare peak time within the

uncertainty of a few minutes. (From the SEP point of view this is called "prompt impulsive". For

delayed impulsive see below and next section.)

Also, the spectral shapes of the in situ observed electron spectrum and hard X-ray photo

spectrum are statistically correlated. This suggests that a common accelerator produces both the

hard X-ray emitting electrons and the escaping electrons (Droege 1996, Krucker et al. 2007).

However, the observed correlation does not agree with the prediction of the thick-target model

(Brown 1971) that would be expected for footpoint emission. This is currently not understood.

Hard X-ray imaging of the solar source region of prompt events often reveals three sources.

Frequently jets are observed in EUV and soft X-rays that appear to emanate from the hard X-ray

sources. This geometry is similar to that of interchange reconnection models, where closed field

lines reconnect with open fields. The jet indicates the direction of open field lines along which

electrons escape. Prompt events are well correlated with the occurrence of 3He rich SEP events.

However, arrival time studies of ions seen at 1 AU suggest a delayed release of the ions

relative to electrons, at least when assuming scatter-free transport. This is rather puzzling as it

suggests a different accelerator for electrons and ions despite the closely connected occurrence.



However, electrons and ions could still be released simultaneously with propagation effects

explaining the observed delays.

This objective is about exploring if the comparative delay is real or due to propagation

effects. It should be studied in different heliocentric distances and hopefully near the perihelion in

order to better establish the nature of the flare region (with STIX & EUI).

Complete set of observations is required: STIX, EUI, RPW + ground-based radio

observations (good to have). That will allow studying the magnetic structures along which flare-

accelerated particles escape into interplanetary space.

3.1.2.6 Delayed events (between X-ray peak and electron release time)

These are events that show sufficiently long delay times (>10 minutes) between the peak

hard X-ray emission and the calculated electron release times.

• Is there indeed a delay between the peak hard X-ray emission and the calculated electron

release times or propagation effects are producing the observed delayed onsets?

• The origin of the delayed release times is currently not understood, but two main ideas are

discussed:

• The delay could be due to time-extended electron acceleration and/or electron

storage at high coronal altitudes during solar flares in combination with a delayed

access of these electrons to magnetic field lines open to interplanetary space

(Laitinen et al., 2000; Classen et al., 2003; Klein et al., 2005; Aurass et al., 2006).

• The delay could be the result of electron acceleration by coronal shocks that move

away from the flare site (see also next section 3.2) (Krucker et al., 1999; Haggerty

and Roelof, 2002; Simnett et al., 2002). The delayed release would be produced by

the time that the shock takes to form and to efficiently accelerate electrons (Mann et

al. 1995; Warmuth & Mann 2005), and/or by the time it takes for the shock to reach

magnetic field lines that are connected to the spacecraft. Problem: shocks not very

efficient at accelerating electrons at higher (MeV) energies (but it might work for

shocks that further accelerate a previously produced population of energetic

electrons, e.g. Selkowitz & Blackman 2007).

• In both cases, the related hard X-ray emissions are expected to be seen after the main

hard X-ray bursts. They are likely to be much fainter and to originate from a different

location. However, detection of these faint emissions will be difficult in the presence

of decaying thermal X-ray emissions at these later times. The best chance of

detection will be for occulted flares where the main flare emission is not yet visible.

The faint emission related to the delayed release may be seen on the disk or at high

altitude above the disk.

3.1.2.7 How are so many electrons accelerated on such short time scales to explain the

observed hard X-ray fluxes?



Hard X-ray imaging observations show most prominent emissions from footpoints of flare

loops in the chromosphere where the ambient density is high enough to stop flare-accelerated

electrons by collisions (e.g. Hoyng et al. 1981). However, fainter, co-temporal hard X-ray sources

are also seen in the corona (e.g. Frost & Dennis 1971, Masuda et al. 1994, Veronig & Brown 2004,

Battaglia & Benz 2006, Krucker et al. 2007) consistent with electron acceleration in the corona. In

particular, RHESSI observations of partially-occulted flares show that at least 90% of all flares have

coronal hard X-ray sources (Krucker & Lin 2008). Further evidence for a coronal acceleration

region comes from radio observations (e.g. Benz 1985, Aurass et al. 2004, Mann et al. 2006). The

details of the transport of electrons from the coronal acceleration site down to the hard X-ray

footpoints are still unclear (e.g. Miller et al. 1997, Önel et al. 2007, Battaglia & Benz 2007).

It is becoming increasingly clear that propagation of energetic electrons does not follow a

simple collisional thick-target scenario, so more sophisticated models of electron transport are

required including effects of non-uniform plasma ionization (e.g. Kontar et al. 2003), return current

(e.g. Zharkova & Gordovskyy 2006), and beam-plasma interaction via various plasma waves (e.g.

Kontar 2001). The deposited energy of the non-thermal electrons heats the chromospheric plasma

and the resulting overpressure drives the hot plasma up the legs of the magnetic loops (e.g. Brown

1973) in the process termed chromospheric evaporation. Hard X-ray observations provide thermal

diagnostics of the heated flare loops.

Observing plans for this objective should include in particular partially limb-occulted flare

observations and observations of hard X-ray emissions associated with CMEs:

Partially limb-occulted flare observations

Partial limb-occultation will frequently provide view angles from which purely coronal

emission (e.g. Hudson 1978) can be readily seen because the footpoint sources are occulted. This

will allow us to study faint coronal sources that otherwise would have been lost in the limited

dynamic range (≤20) of indirect imaging instruments. Coronal hard X-ray emission in the absence

of hard X-ray footpoint emissions, but with the thermal flare loop still partially visible, is best

studied for flares occurring up to ~20° behind the limb. This suggests that for a single spacecraft

about 20% of the observed flares have occulted footpoints with the main flare loop still partially

visible above the limb. In at least 90% of all such events, non-thermal coronal emission is observed,

most prominently during the rise of the thermal emission (Krucker & Lin 2008). Most often the

non-thermal emission is seen close to the thermal loop, but occasionally from above the thermal

flare loops similar to what was reported for the Masuda flare (Masuda et al. 1994). As electron

acceleration is thought to occur in the corona, hard X-ray imaging spectroscopy provides crucial

information on the location and spectrum of energetic electrons before they precipitate to the

chromosphere. STIX will provide frequent partially limb-occulted observations at very high

sensitivity and will be able to see up to 15 weaker coronal emissions than RHESSI. This will allow

us to image hard X-ray emission from the corona at the highest ever sensitivity, free from the

intense footpoint sources, thus providing unique information about the suprathermal electrons

closest to the site in the corona where their acceleration is believed to occur.

Coronal phenomena in hard X-rays associated with CMEs



The partial occultation of a solar flare by the solar limb is an excellent opportunity for

studying faint coronal HXR emissions without competition from the very bright emission of the

footpoint sources. For flares occurring >25° behind the solar limb (i.e. occultation height larger than

0.1 of a solar radius), not are only the hard X-ray footpoints occulted, but also the main thermal and

non-thermal emissions from the corona. Nevertheless, these highly occulted events also show hard

X-ray emission (e.g. Kane et al. 1992) that is associated with fast backside CMEs (Krucker et al.

2007). The emission is faint but has a rather flat (hard) spectrum indicating that the emission is

produced by non-thermal bremsstrahlung of energetic electrons. Multi-spacecraft observations

reveal that the onset of the emission is simultaneous with the onset of the main hard X-ray emission

seen in footpoints (e.g. Kane et al. 1992), but has a much larger source size that expands with time.

These sources move rapidly (~1,000 km s-1) upwards (Hudson et al. 2001, Krucker et al. 2007) in

the same direction as the associated CME. The high coronal emissions may be produced by flare-

accelerated energetic (>10 keV) electrons trapped in magnetic structures related to the CME, or they

may be accelerated in CME current sheets or other coronal magnetic restructuring related to the

CME. However, the details are not understood. The relative number of non-thermal electrons is

observed to be about 10% of the number of thermal electrons in the high coronal source and the

pressure exerted by the non-thermal electrons may, therefore, be comparable to that of the thermal

plasma itself. High-altitude coronal hard X-ray sources are believed to be a common phenomenon

since RHESSI detected them in association with all fast (>1500 km/s) backside CMEs with flare

locations between ~25° and ~50° behind the limb (Krucker et al. 2007). However, present-day

observations are only sensitive to high coronal emissions related to large X-class flares. With

the 15-times enhanced sensitivity of STIX when near the Sun, high coronal emission from backside

CMEs related to M-class flares will regularly be detected.

3.1.2.8 Explore the type III radio bursts delays

Impulsive SEP events are generally accompanied by solar type III radio bursts, indicative of

electrons escaping from the Sun, that drift down to near the local plasma frequency (tens of kHz).

At 1 AU, for most impulsive events the injection of the >25 keV electrons at the Sun, inferred from

the observed velocity dispersion, is delayed by ~10-30 min after the type III radio bursts (Krucker et

al., 1999; Haggerty & Roelof, 2002). These delays have been suggested to be due to propagation

effects in the interplanetary medium (Cane & Erickson 2003; Cane 2003), or to delayed acceleration

by large-scale coronal transient (EIT or Moreton) waves or by shock waves associated with narrow

CMEs (Simnett et al. 2002; Rouillard et al., 2012), or by coronal magnetic restructuring in the

aftermath of CMEs (Maia & Pick 2004; Klein et al. 2005; van Driel-Gesztelyi et al., 2014). Recent

analyses of highly scatter-free impulsive electron events show evidence for two injections (Wang et

al. 2006): a low-energy (~0.4 to ~12 keV) injection that begins ~10 min earlier than the type III

radio burst, and a high-energy (>~13 to 300 keV) injection that starts ~10 min after. This confirms

that type III bursts are produced by ~1-12 keV electrons, consistent with the type III-producing

Langmuir waves being detected in situ at the time of their arrival at 1 AU. Recent theoretical studies

suggest that some of the delays for higher energy electrons may be due to propagation effects.

In a few events, the injection of energetic ions can be inferred, and they lag behind the

electrons by ~30-60 minutes more. SOHO EIT and LASCO images suggest an association of

impulsive events with jets or narrow fast CMEs (Wang et al 2006); perhaps they

accelerate electrons lower down and ions higher up in the corona. However, the pattern of high



charge states observed in impulsive events suggests the ions are accelerated in relatively high-

density regions.

3.1.3 Relativistic electron acceleration (including sub-objectives 3.1.3.1 & 3.1.3.2)

Relativistic electrons were observed up to tens of MeV energies in the 1970s and 80s

(Datlowe 1971, Moses et al. 1989), but in the past 30 years, there have been no measurements at

those high energies to compare with the much more sophisticated observations of ions and

electromagnetic emission. In contrast to keV - tens of keV electrons, relativistic electrons generally

exhibit a diffusive flux-time profile, indicating that substantial scattering has occurred in their

propagation to ~1 AU. Electrons are as fundamental as ions to the understanding of the energy

release process. Their spectral shapes provide a nearly perfect diagnostic of the SEP event type

(Moses et al. 1989) - gradual events show single power laws in momentum, while impulsive events

all show a spectral hardening starting at around 5 MeV/c. Although the hardening of the spectrum is

universally observed, the spectral indices vary from event to event. Another surprising result (from

combined Helios/ISEE-3 electron spectra) is that the spectral shapes are apparently invariant as a

function of the azimuthal distance to the flare if the fluxes are adjusted for different radial distances.

Little is known as to the origin of these electrons or the reason for the hardening. The double power

law spectra have been discussed in terms of a superposition of two electron populations, one

accelerated in flaring loops by a stochastic mechanism, and the other by a shock in the high corona

(Dröge 1996a). How shocks can accelerate electrons to relativistic energies (never observed for

shocks near 1 AU) near the Sun is still an unsolved mystery: can coronal shocks accelerate electrons

to relativistic energies starting from a quasi-thermal population, or is a more energetic seed

population necessary?

3.1.4 Other high-sensitivity X-ray studies

3.1.4.1 Hard X-ray emission of escaping electron beams (thin-target emission)

Electron beams escaping into interplanetary space produce faint hard X-ray emission along

their coronal path in the so-called thin-target approximation at a level several orders of magnitude

smaller than the main flare emissions. The only chance of observing this emission is in events that

are well over the limb such that both the hard X-ray footpoints and the thermal X-ray sources are

occulted. Theoretical calculations show that, under extremely favorable conditions, RHESSI

observations could have enough sensitivity to detect thin-target emission from escaping electrons

(Saint-Hilaire et al. 2008). Hard X-ray emission temporally correlated with radio type III bursts was

observed by RHESSI from an elongated hard X-ray source in the corona, possibly outlining the

electron escape path. However, the emission is about an order of magnitude too bright for purely

thin-target emission from the number of escaping electrons seen near 1 AU by WIND/3DP. STIX

will also not be able to regularly observe thin-target emission from escaping electrons, but, under

favorable conditions such as a hard electron spectrum and high ambient plasma density, STIX will

provide the first clean detections of purely thin-target emission from escaping electron beams.

3.1.4.2 X-ray emission from electrons accelerated at CME shocks



Solar Orbiter will also make it possible to search for hard X-ray emission from electron

beams accelerated at the CME shocks that produce type II radio bursts. As the density of the

emission region of a type II burst is low (<10^9 cm-3), the related hard X-ray emission is expected

to be very faint and only upper limits have been derived so far (Klein et al. 2003). In any case, the

detection will only be possible in highly occulted flares. The high sensitivity of STIX down to low

energies (~4 keV) together with the large number of partially occulted flares with type II bursts will

greatly help to find an event or at least provide more stringent upper limits. Such observations

would identify where along the shock electrons are accelerated, and would provide the first

quantitative measurements of current unknowns, such as the spectrum and energy content of the

electrons associated with radio type II bursts (e.g. Mann & Klassen 2005). Because there have never

been hard X-ray observations of these shock fronts, these observations could provide decisive tests

for theoretical models (e.g. Cairns et al. 2003).

3.2.3.2 3.2 How are energetic particles released from their sources and distributed in space

and time?


SEPs associated with CME-driven shocks have been long known to often arrive at Earth

orbit hours later than would be expected based on their velocities. There are two alternate processes

that might cause this. (1) The acceleration may require significant time to energize the particles

since they must repeatedly collide with the shock to gain energy in many small steps, so the process

may continue for many hours as the shock moves well into the inner solar system. Or (2) the particle

intensities near the shock may create strong turbulence that traps the particles in the vicinity of the

shock, and their intensity observed at Earth orbit depends on the physics of the particles escaping

from the trapping region. Once free of the vicinity of the shock, SEPs may spiral relatively freely on

their way to earth orbit, or more usually will be scattered repeatedly from kinks in the IMF, delaying

their arrival further. The amount of scattering in the interplanetary space varies depending on other

activity such as recent passage of other shocks or solar wind stream interactions. By the time the

particles reach Earth orbit, they are so thoroughly mixed that these effects cannot be

untangled (Gopalswamy et al. 2006; Cohen et al. 2007).

Particles accelerated on magnetic loops can reach very high energies in seconds after the

onset of flaring activity, and then collide with the solar surface where they emit gamma radiation.

There is a poor correlation between the intensity of the gamma radiation and the SEP intensities

observed at Earth orbit, so most particles from this powerful acceleration process do not escape.

Much more common are flare events observed in UV and X-rays that produce a sudden acceleration

of electrons. The electrons can escape from the corona, producing nonthermal radio emission as

they interact with the local plasma. Moving from higher to lower frequencies as the local plasma

density decreases with altitude, the (type III) radio emission makes it possible to track the energetic



electron burst into interplanetary space where it may pass by the observer. Energetic ions, greatly

enriched in 3He and heavy nuclei, accompany these electron bursts (Lin 2006; Mason 2007).

Key open questions in shock associated events are whether particle arrival delays at 1 AU

are due to the length of time needed to accelerate the particles, or due to trapping in the turbulence

near an accelerating shock or a combination of both? For particles accelerated on loops, are the

electrons and ions accelerated from sites low in the corona or at higher altitudes, and how are they

related to the X- and gamma-ray signatures?


Solar Orbiter will revolutionize our understanding of SEP acceleration associated with

CME-driven shocks by probing the inner heliospheric sites where particle acceleration and release

take place. Solar Orbiter will observe how shocks evolve, and whether they are still accelerating

particles as they pass by the spacecraft. If particle arrivals are controlled by the time it takes the

shock to accelerate them, then the highest energy particles will be delayed since they require many

more interactions with the shock. If trapping and release control the timing, then as the shock moves

by the faster and slower particles will have similar intensity changes. Since Solar Orbiter will

simultaneously measure the turbulence properties in the shock acceleration region, it will be

possible to construct a complete theory and models of the acceleration process, and its radial

dependence in the inner heliosphere.

For SEPs accelerated on loops or in reconnection regions, Solar Orbiter will see the coronal

location from UV and X-rays, and then trace the progress of released electrons by radio emission

that will drift to the plasma frequency at the spacecraft for those bursts that pass by. This

unambiguously establishes that the magnetic field line at Solar Orbiter connects to the coronal UV

and X-ray emission site. Since Solar Orbiter can be connected to active regions for periods of days,

this will provide multiple tracings between the heliospheric magnetic field and its origin in the

corona. The corotation phase of Solar Orbiter will considerably lengthen the periods of connection

to active regions, greatly increasing the number of field line origin sites that can be determined from

a single active region. X-ray emission from the flaring sites can be used to derive the energetic

electron spectrum at the flare site, which in turn can be compared with the escaping population to

see if most of the accelerated electrons are released (usually most do not escape). Thanks to the

1/r^2 intensity advantage, Solar Orbiter will observe thousands of these cases and thereby permit

detailed mapping of coronal sources and the trapping properties of the acceleration sites.

3.2.0 What controls the escape of the particles to the heliosphere?

For gradual SEP events, the self-generated turbulence around CME-driven shocks may be

important. Theoretical modeling has pointed out a few possible mechanisms: the particles escape

before large wave growth occurs, adiabatic focusing in the diverging magnetic field, and/or the

shock arriving at a region less favorable to wave growth. On the other hand, if strong turbulence

exists in the corona (i.e., not self-generated), particle acceleration and escape are likely to be

consecutive processes. For impulsive SEP events, self-generated turbulence is likely to be weak, and

escape will depend on the SEPs reaching magnetic field lines connected to the heliosphere. Timing,

compositional, spectral and anisotropy observations for SEPs, together with in situ wave and



magnetic field measurements close to the Sun and field line tracking are key to discriminating

between escape mechanisms.

3.2.1 How do energetic particles scatter and move along the interplanetary magnetic field?

(including sub-objectives 3.2.1.1, 3.2.1.2 & 3.2.1.3)

After SEPs leave their acceleration sites - flares or CME-driven shocks or high coronal

structures - they propagate along the open coronal magnetic fields that expand rapidly close to the

Sun, and then along the large spatial scale interplanetary magnetic field (IMF) - Parker's spiral for

steady flows but highly complex in transients such as ICMEs. The IMF appears to be random and

turbulent at smaller spatial scale. Understanding how energetic particles propagate from their

sources to the point of observation is not only essential to better understand the acceleration

processes on or near the sun, but an important problem in its own right. Propagation of charged

particles parallel to the IMF is affected by two competing processes: the adiabatic motion due to the

divergence of the large-scale IMF results in “focusing” wherein particle pitch angles decrease on

average as they move out; and scattering in pitch angle by small-scale magnetic irregularities at a

rate depending on the strength of the turbulent magnetic field.

A crucial parameter to solve the transport equation is the Fokker-Planck coefficient, Duu,

which can be related to the particle's mean free path if the power spectrum of the turbulence is

known. To understand SEP observations at 1 AU, model calculations of the transport equations are

fit to observation, with assumptions on the r-dependence of the turbulent magnetic field along the

Parker spiral from the acceleration site to 1 AU, since we can only obtain wave spectra at 1

AU. Solar Orbiter’s in situ magnetic field and plasma measurements will map the power spectrum

of the turbulent magnetic field as a function of heliocentric distance in to 0.28 AU.

Transport of ions and electrons may be quite different because electrons and ions resonate in

different regions of the turbulence power spectrum, leading to a rigidity dependence of the particle's

mean free path. Whereas high-rigidity particles are suited to probe the geometry of the fluctuations,

low-rigidity particles are sensitive to the dissipation range and dynamical and thermal effects. Some

models (Dröge, 2003), using power spectra at 1 AU, are able to correctly account for the

dependence of the scattering mean free path on the particle's rigidity and have significantly

improved the agreement between scattering theory and observations. It is important to examine if a

similar rigidity dependence also holds at 0.28 AU.

3.2.2 Latitudinal and longitudinal transport of SEPs

The unique perspective of Solar Orbiter from its inclined orbit in the high latitude phase

should aid in the study of the latitude distribution of gradual and impulsive SEP events close to the

Sun. Ulysses measurements in the outer heliosphere (>1 AU) show that the difference in latitude

between Ulysses and the associated flare orders both onset times and times to maximum well (Dalla

et al. 2003). Sanderson et al. (2003) analyzed the same set of events and found that SEPs either



propagate at the highest heliographic latitudes along magnetic field lines and not across them at

Ulysses’ position. Since onset delay is not correlated with CME parameters, these authors and

Struminsky et al. (2006), conclude that the SEPs either propagate along the distorted magnetic field

with large latitudinal excursions (e.g. Fisk 1996) or diffuse across magnetic field lines close to the

Sun to reach high latitudes. Wibberenz and Cane (2006) suggested lateral coronal transport via low

coronal magnetic loops to explain multi-spacecraft observations including Helios. Thus, high

latitude measurements of SEPs, plus field line tracking, will be crucial to understanding lateral

transport.

We, therefore, need to study the following:

• Propagation at high heliographic latitudes along magnetic field lines vs across them

(Sanderson et al., 2003).

• Propagation along the distorted magnetic field with large latitudinal excursions

(Struminsky et al., 2006).

• Diffusion across magnetic field lines close to the Sun to reach high latitudes

(Dressing et al., 2014).

• Lateral coronal transport via low coronal magnetic loops (Wibberenz and Cane,

2006; Klein et al., 2008).

• Wide angular spread of SEP events. Relative role of cross-field diffusion and shocks

(Dresing et al., 2014; Lario et al., 2014; Gomez-Herrero et al., 2015).

3.2.3 Properties and distribution of near-Sun shocks, their fluctuations and particle

acceleration

The answers to fundamental questions about the sources and acceleration mechanisms of the so-

called large gradual SEPs, whether flares or shocks, depend on many factors including magnetic

connection to active regions that produce flares (e.g., Cane et al., 2006) and on properties of the

CME-driven shocks in the solar corona and the inner heliosphere (Tylka et al. 2005; Desai et al.,

2006a). We need to understand the coronal conditions under which the shocks form and determine

the interplanetary conditions where they evolve. We must then use this information to characterize

shock structure, determine whether they are locally quasi-perpendicular or quasi-parallel, and then

study kinetic properties of shock-related waves and turbulence that control ion scattering mean free

paths near the shock. With geometric and microphysical shock properties determined, we can infer

the rates of particle acceleration and the injection energies as a function of distance and along

different parts of the shock. We need to derive particle acceleration rates near shocks, and by

measuring shocks at many different radii, we will determine the acceleration rates as a function of

distance from the Sun.

• Detect shocks and determine their fine-scale structure (in situ trigger burst mode) at a range

of Mach numbers (Cane and Lario, 2006).



• Measure the waves and turbulence around the shocks and compare these with particle

observations and models in order to determine the acceleration processes of the particles

(Sandroos and Vainio, 2006).

3.2.4 How do large and small-scale structures modulate particle fluxes?

The sharp decreases of high energy cosmic ray fluxes at the Earth, known as Forbush

decreases (e.g. Ifedili, 2004), are caused the passage of “magnetic barriers” in the solar wind. As

well as magnetic clouds, these barriers are often the sheaths of compressed solar wind ahead of

ICMEs, or compressed CIRs (Clack et al., 2000). Such compressions lead to “planar magnetic

structures” (e.g. Jones et al., 2002: see Figure below) and it appears to be these sheets of magnetic

field that efficiently block the particles (Intriligator et al., 2001). However, the efficiency of these

barriers, as they develop close to the Sun, is not known.

• Study of the planarity and large-scale structures of the magnetic field within and around

CIRs and ICMEs and investigate their effect on the fluxes of solar energetic particles and

cosmic rays as these structures evolve with heliocentric distance.

At a much smaller scale, “particle channels” of dramatically enhanced or reduced particle flux

during solar particle events, lasting only around an hour, as well as comparable duration burst of

Jovian electrons observed several AU from the planet, demonstrate the existence of very small scale

magnetic connections and disconnections from particle sources. Solar Orbiter will travel much

closer to the solar sources of particles and, by nearly co-rotating with the Sun, distinguish temporal

variability from spatial structures.

• Determine the small-scale diffusion of particles around and within particle channels, as well

as the spatial and temporal scales of these structures and their connectivity to the Sun.

3.2.5 Shock-surfing acceleration mechanism

See Objective 2.3.4.

3.2.6 Effects of energetic particles propagating downward in the chromosphere

• Generation of heating, shocks, sunquakes in the underlying photosphere (Martinez-Oliveros

et al., 2007).

This objective does not appear in any SOOP. Should we include it

in R_SMALL_HRES_HCAD_WaveStereoscopy

and/or R_SMALL_HRES_HCAD_Photospheric_Dynamics_Structure?

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/R_SMALL_HRES_HCAD_WaveStereoscopy




3.2.3.3 3.3 What are the seed populations for energetic particles?


The low-energy particles accelerated by CME-driven shocks to SEP energies are called the

seed population. The observed ionization states of SEP ions show temperatures typical of the

corona, ruling out hot material on flare loops as the seeds. But SEPs also show significant

abundances of ions such as 3He and singly ionized He, which are virtually absent from the solar

wind. The observed energetic particle abundances indicate that the suprathermal ion pool, composed

of ions from a few to a few 10 s the speed of the solar wind, is the likely source. At 1 AU, the

suprathermal ion pool is ~100 times more variable in intensity than the solar wind and varies in

composition depending on solar and interplanetary activity. The suprathermal ions are continuously

present at 1 AU, but it is not known if there is a continuous solar source, or if these ions are from

other activities such as acceleration in association with fast and slow solar wind streams. Inside 1

AU, the suprathermal ion pool is expected to show significant radial dependence due to the different

processes that contribute to the mixture, but it is unexplored (Desai et al. 2006; Mewaldt et al. 2007;

Lee 2007; Fisk and Gloeckler 2007).

For SEPs accelerated on loops or in reconnection regions that give rise to electron and type-

III radio bursts, ionization states are coronal-like at lower energies and change over to much hotter

flare-like at high energies. This may be evidence for a complex source, or, more likely, of energetic

particle stripping as the ions escape from a low coronal source. For SEPs accelerated at

reconnection sites behind CMEs abundances and ionization states would be coronal (Klecker et al.

2006).

Critical questions in this area are: what is the suprathermal ion pool in the inner heliosphere,

including its composition and temporal and spatial variations? What turbulence or stochastic

mechanisms in the inner heliosphere accelerate particles to suprathermal energies? Are the source

locations and arrival times of electrons from SEPs on loops or reconnection regions consistent with

a low or high coronal source?


By systematically mapping the suprathermal ion pool in the inner heliosphere with

spectroscopic and in-situ data, Solar Orbiter will provide the missing seed particle data for models

of SEP acceleration associated with shocks. Together with the shock and turbulence parameters also

measured on Solar Orbiter, there will be the first well-constrained models. Since the suprathermal

ion pool composition varies, different shock events will be expected to produce correspondingly

different energetic particle populations that can be examined on a case-by-case basis. The high-

latitude phase of the mission will add an important third dimension to the suprathermal pool

mapping, since it will be more heavily influenced by, e.g., mid-latitude streamer belts, making it

possible to probe the solar and interplanetary origins of the seed particle populations. Taken

together, these observations will make it possible to construct the first complete physics-based

theory and models of particle acceleration close to the Sun.



For SEPs accelerated on loops or in reconnection regions the 1/r^2 advantage of Solar

Orbiter will again provide a decisive advantage since particle properties will be accurately measured

and compared with much more precise information on the coronal location. This will permit

distinguishing between low coronal sources that result in stripping of escaping particles vs. higher

sources which could mimic stripping properties. SEPs accelerated from reconnection regions in the

back of CME lift-offs will be identified by comparing energetic particle timing with the location of

the CME, and energetic particle composition with that determined spectroscopically for the remote

coronal source.

3.3.1 What are the properties and distribution of suprathermal seed populations?

The suprathermal population contains a mix of solar wind ions, flare populations, gradual

SEP events, pickup ions, and other particle populations (see figure). Since the populations in this

mix have distinct histories, suprathermal elemental and charge-state composition are likely to vary

with energy (Fisk and Gloeckler 2007), time (Desai et al., 2006a, b) and location. The high initial

speeds of suprathermal ions predispose them to efficient injection into acceleration near shocks,

making them the ideal seed population for SEPs. Where shocks are quasi-perpendicular, we expect

injection energies to be high, and the SEP population should resemble the composition of the higher

energy suprathermal ions. Where shocks are quasi-parallel, we expect the SEP composition to

reflect the lower energy portions of the suprathermal population.

3.3.1.1 Characterization of the suprathermal population

▪ Solar wind ions

▪ Flare populations

▪ Gradual SEP events

▪ Pickup ions

▪ Other particle populations

3.3.1.2 Characterize the suprathermal elemental and charge-state composition as a function of

energy, time, and location.

Characterize the suprathermal elemental and charge-state composition as a function of energy (Fisk

and Gloeckler, 2007), time (Desai et al., 2006a, b), and location.

3.3.1.3 Role of shocks in generating SEPs

• Role of shocks in generating SEPs: comparison between the composition of suprathermal

ions near the shocks and that of energetic particles.

• Distinction for quasi-perpendicular and quasi-parallel shocks.



• Spectroscopic observations of CME shocks’ line profiles.

• Visible light and UV observations of shocks in the inner corona.



3.2.4 Objective 4: How does the solar dynamo work and drive connections between the

Sun and the heliosphere?

The Sun’s magnetic field dominates the solar atmosphere. It structures the coronal plasma,

drives much of the coronal dynamics, and produces all the observed energetic phenomena. One of

the most striking features of solar magnetism is its ~11-year activity cycle, which is manifest in all

the associated solar and heliospheric phenomena. Similar activity cycles are also observed in a

broad range of stars in the right half of the Hertzsprung-Russell diagram, and the Sun is an

important test case for dynamo models of stellar activity.

The Sun’s global magnetic field is generated by a dynamo generally believed to be seated in

the tachocline, the shear layer at the base of the convection zone. According to flux-transport

dynamo models (e.g., Dikpati and Gilman 2008), meridional circulation, and other near-surface

flows transport magnetic flux from decaying active regions to the poles. There, subduction carries it

to the tachocline to be reprocessed for the next cycle. This ‘conveyor belt’ scenario provides a

natural explanation for the sunspot cycle and characterizing the flows that drive it will provide a

crucial test of our models and may also allow us to predict the length and amplitude of future cycles.

However, current models fail miserably at predicting actual global solar behavior. For example, the

current sunspot minimum has been far deeper and longer than predicted by any solar modeling

group, indicating that crucial elements are missing from current understanding.

A major weakness of current global dynamo models is the poor constraint of the meridional

circulation at high latitudes. The exact profile and nature of the turnover from poleward flow to

subduction strongly affect behavior of the resulting global dynamo (e.g., Dikpati and Charbonneau

1999), but detecting and characterizing the solar flow is essentially impossible at shallow viewing

angles in the ecliptic plane.

In addition to the global dynamo, turbulent convection may drive a local dynamo that could

be responsible for generating the observed weak, small-scale internetwork field, which is ubiquitous

across the surface and appears to dominate the emergent unsigned flux there.

A key objective of the Solar Orbiter mission is to measure and characterize the flows that

transport the solar magnetic fields: complex near-surface flows, the meridional flow, and the

differential rotation at all latitudes and radii. Of particular and perhaps paramount importance for

advancing our understanding of the solar dynamo and the polarity reversal of the global magnetic

field is a detailed knowledge of magnetic flux transport near the poles. Hinode, peering over the

Sun’s limb from a heliographic latitude of 7°, has provided a tantalizing glimpse of the Sun’s high-

latitude region above 70°; however, observations from near the ecliptic lack the detail, coverage,

and unambiguous interpretation needed to understand the properties and dynamics of the polar

region. Thus, Solar Orbiter’s imaging of the properties and dynamics of the polar region during the

out-of-the- ecliptic phase of the mission (reaching heliographic latitudes of 25° during the nominal

mission and as high as 34° during the extended mission) will provide urgently needed constraints on

our models of the solar dynamo.

Most of the open magnetic flux that extends into the heliosphere originates from the Sun’s

polar regions, from polar coronal holes. The current solar minimum activity period, which is deeper

and more extended than previously measured minima, demonstrates the importance of this polar

field to the solar wind and heliosphere.



There is evidence that the solar wind dynamic pressure, composition and turbulence levels,

as well as the strength of the heliospheric magnetic field, have all changed in the last few years in

ways that are unprecedented in the space age. None of these changes were predicted, and current

solar conditions present a challenge to our understanding of the solar dynamo and its effects on the

solar system at large and the Earth in particular.

In the following sections, we discuss in more detail three interrelated questions that flow

down from this top-level question: How is magnetic flux transported to and reprocessed at high

solar latitudes? What are the properties of the magnetic field at high solar latitudes? Are there

separate dynamo processes in the Sun?

3.2.4.1 4.0 Overall remarks and feasibility concerning Objective 4 observations with Solar

Orbiter

The most important assets of Solar Orbiter/PHI for helioseismology are:

• Combine Earth-based, front-side observations with PHI observations from back-side (and if

possible higher latitudes) to

o be able to observe front and back side sun simultaneously for few days, which would

strongly improve local helioseismology

o improve and calibrate the far side modeling based on helioseismology

o Use PHI’s out-of-ecliptic observations to measure meridional flow at high latitudes

For local helioseismology, PHI would need to observe at 1 min cadence and good resolution. Only

Dopplergrams are needed, i.e. 1 data product out of the 5 in PHI’s standard dataset.

Helioseismology observing strategies & remarks

• The requirements for local helioseismology and feature tracking are a moderate spatial

resolution (f-mode wavelength; a few Mm) with a temporal cadence of 60s.

• For granulation feature tracking, one needs to resolve granules. The cadence should

be 60s between a pair of observations, then a waiting time of 30 min can be introduced

before recording the second pair of two 60s separated images. Onboard processing is

possible to obtain the flow map.

• For supergranulation tracking, images can be taken separated by 1h.

Again, onboard processing is possible. Granulation needs to be suppressed, which requires

averaging on board. Furthermore, the solar rotation needs to be taken into account. High-

resolution images are best taken close to the disk center to minimize limb effects.

• Noise reduction requires long time series. Given the noise level that increases with latitude

and the smaller number of features at the poles, the error on the horizontal flow speed goes

up to >6m/s at the poles with a time series of 30 days length, i.e. the meridional flow might

not be measurable at the poles.

• Compression can help reduce the data rates with little effect on time-distance

helioseismology (JPEG) and local correlation tracking (quantization).



• Studying the center-to-limb effect in local helioseismology could help understanding the

physics of solar oscillation modes.

• Different observing strategies are suggested in order to address the following science goals:

o Near-surface differential rotation: for measuring differential rotation at high latitudes

one set of observations would be helpful. This requires 30 days of data. Resolving

the torsional oscillations requires repetition of measurements. These 30 days of data

do not need to be recorded at once. Averaging of single power spectra obtained from

observing several sequences of 10 days is possible. Both HRT and FDT could be

used here: HRT is only necessary for granulation tracking but the two other

techniques could be based on FDT data.

o Near-surface meridional flow at high latitudes requires 120 days of data due to the

low amplitude of the flow. Again, averaging of data is possible, but this smears out

cycle effects. Observations for this sub-objective make sense in the later phases of

the mission as a minimum latitude of 15-20 degrees is required to start analyzing the

poles (depending on B angle).

o For studying convection at high latitudes, 7 days of data are needed for obtaining

flow maps and useful power spectra. Weeks to months of data would be needed for

statistical analyses.

o Large-scale convective flows

• A possible orbit for all these measurements is MTP11 - 2024/01/01 - 2024/07/01, with Solar

Orbiter at 45 degrees from Earth.

• Furthermore, observations need to be done with FDT or HRT with 60 days of observation

time: FDT at high latitude far from the Sun (0.7 AU) and HRT at high latitudes close to the

Sun (0.4 AU).

• In summary, observing times are the biggest constraint. As a high cadence of 60s is required

for many science goals, the feasibility of the observing programs highly depend on the

compression that can be applied. For each of the science objectives above, the best methods

need to be defined as well as the most suitable compression.

• Some of the helioseismology goals can be addressed with synoptic observations. PHI

synoptic observations outside the RS windows could be helpful but this is not included in the

current baseline of the mission operations.

Measurements and topology of the polar magnetic field

• PHI is optimally used for measuring the solar polar magnetic field at maximum solar latitude

and at minimum distance.

• Co-observations should be done from Earth, i.e. at a large B0 angle in March and September

when Solar Orbiter is observing the pole visible to Earth. September is preferable as

observations are then possible from the solar telescopes on the Canary Islands. In the current

planning (October 2018 Option E), a suitable orbit for this is MTP14 - 2025/07/01 -

2026/01/01.

• These observations should be combined with other observables, especially EUI observing

polar jets and SPICE.

• Polar magnetic field useful publications include Tsuneta et al., 2008, ApJ 688, 1374 and

Shiota et al., 2012, ApJ, 753, 157.

https://issues.cosmos.esa.int/solarorbiterwiki/pages/viewpage.action?pageId=15075161





Magnetoconvection

• Magnetoconvection is present on the Quiet Sun and in active regions.

• The challenges are the various spatial scales that range from 100 km to 50 Mm, as well as

the various temporal scales ranging from 1 min (small magnetic elements) to 1 week (active

regions).

• Observations from the ground are affected by solar rotation, i.e. the maximum observing

time is 14 days. In addition, projection effects influence the data.

• Zeeman measurements exhibit 180° ambiguity in the azimuthal magnetic field component,

which can be resolved by stereoscopic observations.

• For this objective, it is important to use the good spatial resolution (pixel size of 110 km)

and the large field of view 1000’’x1000’’ of PHI/HRT for combined observations between

Earth-bases telescopes and Solar Orbiter near the almost co-rotating phases to allow long

periods of AR tracking. This should enable studying the long-term behavior of active

regions. The advantage of the almost co-rotation will allow observing many flares and

follow the decay of active regions.

• Solar Orbiter's most important assets for magnetoconvection are:

o longer observations of the same target due to near co-rotation phase close to the Sun,

o Solar Orbiter is an observatory combining many remote sensing instruments,

o combined with ground-based support or SDO, it will provide a new vantage point

and allow stereoscopy.

• It is suggested to use an observing program of 15 days during maximum and declining solar

cycle 25 to follow up AR dynamics.

• An optimal orbit is MTP-07 (0.3 AU, 8-11.5 degree/day, latitude change of 15 degrees).

MTP-10 would be fine too.

• Another observing sequence should include the full remote-sensing package (PHI, EUI,

SPICE) coordinated with DKIST and GREGOR. Observing targets should be quiet sun

regions at the disc center, the limbs and polar regions. Possible orbits are the first orbit and

later orbits of the mission with higher latitudes. The length of the time series should be 3h

per pointing (telemetry intensive observations).

Other related goals

• Measure solar oblateness and luminosity of the Sun. For this, rolls of the spacecraft are

required, with a minimum of 8 positions, preferably close to the equatorial plane. This could

be done during communication or instrument calibration spacecraft rolls.

• Concerning the question whether a small-scale dynamo is acting on the Sun, the difficulty

for observations is that the very convolved field is hard to resolve spatially. Solar Orbiter

could look at the distribution of the magnetic field and the emergence rate per feature over

latitudes. If emergence rate is independent of latitude, probably the feature is produced by

local dynamo rather than local one.

• Regarding the solar irradiance, Solar Orbiter has no measurements of the Sun's luminosity.

However, it can help obtaining synoptic charts with extended viewing windows. Those could

be used for modes of the total solar irradiance.



• Furthermore, Solar Orbiter can help answer the question whether the observations of the

Sun's chromospheric and photospheric activity is different from that observed from other

stars. As one sees the Sun only from the ecliptic, comparisons with other stars might be

biased, since they are observed from all viewing angles. By leaving the ecliptic, we might be

able to test this hypothesis.

Related SOOPs

• The following SOOPs are defined for helioseismology observations:

o R_SMALL_HRES_HCAD_ModePhysics: for most helioseismology objectives.

o R_SMALL_HRES_HCAD_Photospheric_Dynamics_Structure: for addressing flux

emergence in the quiet sun, as well as polar features.

o R_SMALL_MRES_MCAD_AR_LongTerm: for the decay of Active Regions

o R_SMALL_HRES_HCAD_PDF_Mosaic: for the measurement of the probability

distribution function.

3.2.4.2 4.1 How is magnetic flux transported to and re-processed at high solar latitudes?


In the last decades, the mapping of surface and subsurface flow fields at low and middle

latitudes has seen major advances, largely due to the availability of high-quality data from the

SOHO's Michelson Doppler Imager (MDI) instrument. These data have provided accurate

knowledge of differential rotation, the low latitude, near-surface part of the meridional flows, and

the near-surface torsional oscillations, which are rhythmic changes in the rotation speed that travel

from mid-latitudes both equatorward and poleward (Howe et al. 2006). Local helioseismic

techniques have also reached a level of maturity that allows the three-dimensional structure of the

shallow velocity field beneath the solar surface to be determined.

Despite these advances, progress in understanding the operation of the solar dynamo

depends on how well we understand differential rotation and the meridional flows near the poles of

the Sun. However, because of the lack of out-of-the-ecliptic observations, the near-polar flow fields

remain poorly mapped, as does the differential rotation at high latitudes (see Beck 2000; Thompson

et al. 2003). The meridional flow, in particular, the very foundation of the flux transport dynamo, is

not well characterized above ~50° latitude; it is not even certain that it consists of only one cell in

each hemisphere. The return flow, believed to occur at the base of the convection zone, is entirely

undetermined save for the requirement of mass conservation. All these flows must be better

constrained observationally in order to help solve the puzzle of the solar cycle and to advance our

understanding of the operation of the solar dynamo (and, more broadly, of stellar dynamos

generally).


https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/R_SMALL_HRES_HCAD_ModePhysics


https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/R_SMALL_MRES_MCAD_AR_LongTerm

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/R_SMALL_HRES_HCAD_PDF_Mosaic



Solar Orbiter will measure or infer local and convective flows, rotation, and meridional

circulation in the photosphere and in the subsurface convection zone at all heliographic latitudes

including, during the later stages of the nominal mission, at the critical near-polar latitudes. Solar

Orbiter will reveal the patterns of differential rotation, the geometry of the meridional flow, the

structure of subduction areas around the poles where the solar plasma dives back into the Sun, and

the properties of convection cells below the solar surface. This will be achieved through correlation

tracking of small features, direct imaging of Doppler shifts, and helioseismic observations

(including the first from a high-latitude vantage point). By monitoring the temporal variations over

the course of the mission, it will be possible to deduce solar cycle variations in the flows.

Solar Orbiter will resolve small-scale magnetic features near the poles, even within the

nominal mission phase, and right up to the poles during the extended mission. It will determine the

detailed surface flow field through tracking algorithms. Such algorithms provide only inconclusive

results when applied to polar data obtained from near-Earth orbit due to the foreshortening. Doppler

maps of the line-of-sight velocity component will complement the correlation tracking

measurements and will also reveal convection, rotation, and meridional circulation flows.

Time series of Doppler and intensity maps will be used to probe the three-dimensional mass

flows in the upper layers of the convection zone, at high heliographic latitudes. The flows will be

inferred using the methods of local helioseismology (e.g., Gizon and Birch 2005): time-distance

helioseismology, ring diagram analysis, helioseismic holography, and direct modeling. Using

SOHO/MDI Dopplergrams, it was demonstrated that even complex velocity fields can be derived

with a single day of data (e.g., Jackiewicz et al. 2008).

The deeper layers of the convection zone will be studied using both local and the global

methods of helioseismology. Moreover, Solar Orbiter will provide the first opportunity to

implement the novel technique of stereoscopic helioseismology to probe flows and structural

heterogeneities deep in the convection zone, even reaching down to the tachocline. Combining Solar

Orbiter observations with ground- or space-based helioseismic observations from 1 AU (e.g.,

GONG or SDO) will open new windows into the Sun. Looking at the Sun from two distinct viewing

angles will increase the observed fraction of the Sun’s surface and will benefit global

helioseismology because the modes of oscillation will be easier to disentangle (reduction of spatial

leaks). With stereoscopic helioseismology, new acoustic ray paths can be taken into account to

probe deeper layers in the interior, including the bottom of the convection zone.

4.1.1 Study the detailed solar surface flow patterns in the polar regions, including coronal hole

boundaries.

Solar Orbiter will provide the first detailed view of the polar subsurface layers by carrying out

seismic measurements both from a high-latitude vantage point and "stereoscopically" by combining

PHI data and seismic data from the ground or from NEO (e.g. SDO/HMI). These observations will

reveal the patterns of differential rotation, the geometry of the meridional flow, and the properties of

convection cell below the solar surface. By monitoring the temporal variations over the course of

the mission, it will be possible to deduce solar-cycle variations in the flows.

Thanks to global helioseismology (e.g. Christensen- Dalsgaard 2002), the solar differential rotation

has been mapped as a function of latitude and radial distance throughout most of the convection



zone. For heliographic latitudes above 70°, however, the global oscillation mode inversions are

uncertain, leaving an incomplete picture of the solar interior. Local helioseismology (e.g. Gizon &

Birch 2005) aims to measure the 3D velocity vectors of the material flows in the solar interior,

allowing studies of convective, rotational and meridional flows, as well as sunspots and active

regions. Local helioseismology together with PHI observations will enable the study of high

heliographic latitudes. Another important goal of Solar Orbiter is to implement stereoscopic

helioseismology by combining PHI data with Doppler measurements from Earth or NEO

instruments. Local helioseismic inversions from techniques such as time-distance helioseismology

or helioseismic holography will be able to probe deeper into the Sun using observations from widely

separated vantage points because skip distances of more than half a circumference will, at last,

become accessible. This will be important for probing the tachocline at the base of the convection

zone, where the dynamo is surmised to be situated.

• Track granules and magnetic features to follow their motions and mutual interaction over

time (Abramenko et al., 2011; Giannattasio et al., 2013; 2014; Gosic et al., 2014; Requerey

et al., 2014).

• Track mesogranules and supergranules and determine lifetime, sizes, horizontal velocities,

and other properties such as helicity of the flows (e.g. Hathaway 2000, Palacios 2012).

4.1.2 Study the subtle cancellation effects that lead to the reversal of the dominant polarity at

the poles

Study the subtle cancellation effects that lead to the reversal of the dominant polarity at the poles

(Wang et al., 1989; Sheeley, 1991; Makarov et al., 2003).

4.1.2.1 Studies of convective, rotational and meridional flows.

4.1.2.2 Local helioseismology of polar regions (Gizon and Birch, 2005).

4.1.2.3 MDI-like medium-l program (see Löptien et al., 2014).

4.1.2.4 By which physical processes is magnetic flux at the poles removed?

4.1.2.5 How does the cancellation process at the poles differ from cancellation at lower latitudes?

Are there different rates?

4.1.3 Explore the transport processes of magnetic flux from the activity belts towards the poles

and the interaction of this flux with the already present polar magnetic field

4.1.3.1 Follow the evolution of the magnetic flux at different latitudes over the solar cycle (full

disk).

4.1.3.2 Follow individual magnetic features flux from lower to high latitudes (high spatial

resolution).

4.1.3.3 Determine flow rates (differential rotation, meridional flow) obtained by tracking magnetic

features and compare with those determined from tracking granules and from Doppler shifts.

4.1.3.4 How do supergranular flows facilitate or impede the latitudinal transport of small-scale

magnetic features?

4.1.3.5 Compare measurements of the magnetic flux from high and low heliographic latitudes (full

disk and high resolution). How strongly do (near) Earth-based magnetic field measurements need to

be corrected at higher latitudes?



4.1.4 Study the influence of cancellations at all heights in the atmosphere

4.1.4.1 Understand the origin of polar jets.

3.2.4.3 4.2 What are the properties of the magnetic field at high solar latitudes?


Meridional circulation transports the surface magnetic flux toward the poles, where a

concentration of magnetic flux is expected to occur. However, because of the directional

sensitivity of the Zeeman effect and magnetic polarity cancellation resulting from geometric

foreshortening, present-day observations from the ecliptic at 1 AU can provide only a poor

representation of the polar magnetic field. The high resolution of Hinode’s Solar Optical Telescope

(SOT) can partly overcome the second disadvantage (Tsuneta et al. 2008), but not the first.

Consequently, an accurate quantitative estimate of the polar magnetic field remains a major and as

yet unattained goal.

The polar field is directly related to the dynamo process, presumably as a source of poloidal

field that is wound up by the differential rotation in the shear layer at the base of the convection

zone. The distribution of the magnetic field at the poles drives the formation and evolution of polar

coronal holes, polar plumes, X-ray jets, and other events and structures that characterize the polar

corona. Polar coronal holes have been intensively studied from the non-ideal vantage point offered

by the ecliptic, but never imaged from outside the ecliptic. Consequently, the distribution of the

polar field and the origin of polar structures are only poorly determined. The fast solar wind is

associated with open field lines inside coronal holes, whereas at least parts of the slow solar wind

are thought to emanate from the coronal hole boundaries. Understanding the interaction of open and

closed field lines across these boundaries is of paramount importance for elucidating the connection

between the solar magnetic field and the heliosphere.

The magnetic flux in the heliosphere varies with the solar cycle (Owens et al. 2008). There is

evidence that the heliospheric magnetic flux has increased substantially in the last hundred years,

perhaps by as much as a factor of two (Lockwood et al. 1999; Rouillard et al. 2007), possibly due to

a long-term change in the Sun’s dynamo action. As already noted, however, the interplanetary

magnetic field was dramatically lower than expected during the last solar minimum. Models based

on the injection of flux into the heliosphere by coronal mass ejections cannot explain this reduction,

and it is becoming clear that the processes by which flux is added to and removed from the

heliosphere are more complex than previously thought.


Solar Orbiter’s comprehensive imaging instruments will characterize the properties and

dynamics of the polar regions for the first time, including magnetic fields, plasma flows, and

temperatures. Solar Orbiter will make the first reliable measurements of the amount of polar



magnetic flux, its spatial distribution and its evolution (by comparing results from different orbits),

providing an independent constraint on the strength and direction of the meridional flow near the

pole. The evolution of Solar Orbiter’s orbit to higher heliographic latitudes will make it possible to

study the transport of magnetic flux from the activity belts toward the poles, which drives the

polarity reversal of the global magnetic field (see Wang et al. 1989; Sheeley 1991; Makarov et al.

2003). From its viewpoint outside the ecliptic, Solar Orbiter will probe the cancellation processes

that take place when flux elements of opposite polarity meet as part of the polarity reversal process.

Joint observations from Solar Orbiter and spacecraft in the ecliptic will determine, with high

accuracy, the transversal magnetic field, which is notoriously difficult to measure, along with

derived quantities such as the electric current density.

Solar Orbiter will measure the photospheric magnetic field at the poles, while

simultaneously imaging the coronal and heliospheric structure at visible and EUV wavelengths. In

addition, as the spacecraft passes through the mid-latitude slow/fast wind boundary at around 0.5

AU, the field and plasma properties of the solar wind will be measured. With the help of magnetic

field extrapolation methods these observations will, for the first time, allow the photospheric and

coronal magnetic field in polar coronal holes to be studied simultaneously and the evolution of polar

coronal hole boundaries and other coronal structures to be investigated.

Solar Orbiter’s observations from progressively higher heliographic latitudes (25° by the end

of the nominal mission) will enable the first coordinated investigation (jointly with spacecraft in the

ecliptic) of the three-dimensional structure of the inner heliosphere. These observations will reveal

the links between the Sun’s polar regions and the properties of the solar wind and interplanetary

magnetic field, in particular the heliospheric current sheet, which is used as a proxy for the tilt of the

solar magnetic dipole. In addition, Solar Orbiter will pass both north and south of the solar

equatorial plane in each orbit, with repeated transits through the equatorial streamer belt and

through the slow/fast wind boundary at mid-latitudes into the polar wind, making it possible to

follow the evolution of the solar wind and interplanetary magnetic field as well as of the sources in

the polar coronal holes. Ulysses has shown that poleward of the edge of coronal holes the properties

of the solar wind are relatively uniform, so that Solar Orbiter only needs to reach heliographic

latitudes just above the coronal hole edge to enter the high-speed solar wind. The orbital inclination

of 25° reached during the nominal mission is sufficiently high to satisfy this constraint.

4.2.1 Probability density function (PDF) of solar high-latitude magnetic field structures

4.2.1.1 Study the shape and morphology of solar high-latitude magnetic field structures as well as

their relation to the flow fields.

4.2.1.2 Study the distribution of field strengths (and magnetic vectors) of solar high-latitude

magnetic field structures.

4.2.1.3 Compare the results from the above objective (4.2.1.2) with the field strength distribution of

sunspots and active regions.

4.2.2 Basic properties of solar high-latitude magnetic field structures



4.2.2.1 What is the intensity-field strength relation of polar magnetic field structures and how does it

change with the activity cycle?

4.2.2.2 What is the emergence rate and lifetime of polar magnetic field structures and how do they

change with the activity cycle?

4.2.2.3 Are there differences between the network at polar and at low latitudes?

4.2.3 Probe the structure in deep layers of the Sun

4.2.3.1 Probe the tachocline at the base of the convection zone, where the dynamo is surmised to be

situated.

4.2.3.2 Studies of convective, rotational and meridional flows.

4.2.3.3 MDI-like medium-l program

4.2.3.4 Stereoscopic helioseismology

3.2.4.4 4.3 What is the nature of magnetoconvection?

4.3.1 What are the velocity and magnetic vectors in the solar photosphere?

• Requires simultaneous observations from PHI and an instrument on the ground or another

satellite.

• Simultaneous observations at different viewing angles.

4.3.2 What is the 3D geometry of the solar surface in convective and magnetic features?

• Study the undulating 3D structure of the visible solar surface (Lites et al., 2004).

• Wilson depression appearing in magnetic structures such as faculae, pores and sunspots

(Solanki, 2003).

• Exact estimate of the height difference of the surface layer of different structures.

• Local helioseismology of active regions (Gizon and Birch, 2005).

4.3.3 How does the brightness of magnetic features change over the solar disk?

• Variation across the disk of the contrast of magnetic features relative to the quiet Sun (Topka

et al., 1997; Ortiz et al., 2002; Hirzberger and Wiehr, 2005).

4.3.4 What is the long-term behavior of active regions?

• Study sunspot evolution during co-rotation phase.

• Study the large-scale flows (e.g. sunspot moat flows) in and around active regions.

• Magnetic oscillations in sunspots and active regions.

4.3.5 How do magnetic fields emerge on the solar surface (coalescence, magnetic loops, convective

collapse?).



• See e.g. Lamb et al., 2008; Centeno et al., 2007; Martínez González & Bellot Rubio, 2009;

Nagata et al., 2008.

3.2.4.5 4.4 Are there separate dynamo processes acting in the Sun?


MHD simulations indicate that a local turbulent dynamo should be acting in the Sun’s turbulent

convection zone (Brun et al. 2004) and even in the near-surface layers (Vögler and Schüssler 2007).

Hinode/SOT has detected ubiquitous horizontal magnetic fields in quiet regions of the Sun (Lites et

al. 2007), which are possibly generated by local dynamo action (Pietarila Graham et al. 2009).

These small, weak features (inter-network fields; Zirin 1987) bring 100 times more magnetic flux to

the solar surface than the stronger features that are known to be the product of the global dynamo,

and have themselves shown to be in cross-scale turbulent equilibrium (Schrijver et al. 1997). Even

the smallest observable features have been shown to be formed primarily by aggregation of yet

smaller, yet more prevalent features too small to resolve with current instrumentation (Lamb et al.

2008, 2009). It is, however, still uncertain whether a separate local, turbulent dynamo really is

acting on the Sun and how strongly it contributes to the Sun’s magnetic flux (and magnetic energy).

In particular, all solar magnetic features, from the smallest observable intergranular flux

concentrations to the largest active regions, have been shown (Parnell et al. 2009) to have a power

law (scale free) probability distribution function, suggesting that a single turbulent mechanism may

be responsible for all observable scales of magnetic activity.


One way to distinguish between the products of a global and a local dynamo is to study the

distribution of small elements of freshly emerging magnetic flux over heliographic latitude. The

global dynamo, presumably owing to the structure of the differential rotation and the meridional

flow near the base of the convection zone, leads to the emergence of large bipolar magnetic regions

(active regions) at the solar surface at latitudes between 5° and 30° and of smaller ephemeral active

regions over a larger range of latitudes, but concentrated also at low latitudes. In contrast, a local

turbulent dynamo is expected to enhance field more uniformly across the surface.

Observations carried out from the ecliptic cannot quantitatively determine the latitudinal distribution

of magnetic flux and in particular the emergence of small-scale magnetic features (inter-network

fields) due to foreshortening and the different sensitivity of the Zeeman effect to longitudinal and

transversal fields. Solar Orbiter, by flying to latitudes of 25° and higher above the ecliptic, will be

able to measure weak magnetic features equally well at low and high latitudes (Martínez Pillet

2006). If the number and size (i.e., magnetic flux) distributions of such features are significantly

different at high latitudes, then even the weak features are probably due to the global dynamo. If,

however, they are evenly distributed, then the evidence for a significant role of a local dynamo will

be greatly strengthened. Current work is confounded by viewing angle restrictions near the poles, by

the ubiquitous seething horizontal field (e.g., Harvey et al. 2007), and by small deflections in near-

vertical fields, which dominate observed feature distributions near the limb of the Sun.



Detailed sub-objectives:

4.4.1 Compare the distribution of small-scale fields at low and high latitudes (Tsuneta et al.,

2008; Ito et al., 2010).

4.4.1.1 Is the magnetic network equally strong between low and high latitudes? How does this

distribution change between activity maximum and minimum?

4.4.1.2 What is the latitude distribution of the emergence of ephemeral regions? Is this distribution

dependent on the cycle phase?

The emergence, diffusion and decay of ephemeral regions near the poles and below high-latitude

coronal holes should be studied for the aspect of how they feed the magnetic network (see e.g.

Simon et al. 2001, ApJ 561, 49 427; Gosic et al. 2014, ApJ 797). In particular, the latitudinal

dependence of this decay process would be interesting to study.

4.4.1.3 What is the latitude distribution of internetwork magnetic fields and their emergence rates?

Do they depend on the phase of the solar cycle?

4.4.1.4 What is the latitudinal distribution of the linear-polarization features in the quiet Sun? Is

there a solar cycle dependence?

4.4.2 Joy’s law at high latitudes

4.4.2.1 Extend Joy’s law at low latitudes from large to small bipolar features, i.e. from active

regions to ephemeral regions.

4.4.2.2 Compare the tilt angle distributions of ephemeral regions at low and high latitudes. How do

these distributions change over the cycle?

4.4.3 Understand differences in size and internal structuring of magnetic concentrations in

high and low latitudes and relate them to different origins or dynamic environments (e.g. Lagg

et al., 2010; Martínez González, 2012; Requerey et al., 2014; 2015).

3.2.4.6 4.5 How are coronal and heliospheric phenomena related to the solar dynamo?

4.5.1 Explore the quasi-biennial modulation of galactic cosmic rays (Laurenza et al., 2012) and of

flare-CME onset (Telloni et al., 2015).

4.5.2 Explore the effect of the solar polarity reversal (Gnevyshev gap) on the heliosphere during

solar cycle 25 (Storini et al., 2003) and compare with that of previous cycles.

4.5.3 Explore the possible effect of the Centennial Gleissberg Cycle on the heliosphere during solar

cycles 24 and 25 (Feynman & Ruzmaikin, 2014).

4.5.4 Determine the solar wind, magnetic field, energetic particles and radio emission properties

during solar cycle 25 and compare to those found by Helios.



3.2.5 5. Additional science objectives

This section will be significantly modified: many of the sub-objectives are not additional but instead

part of Objectives 1-4, they have to be reintegrated to the relevant sections.

5.1 Additional Science Objectives of EUI

5.1.1 Study the corona and its phenomena in a high spatial and temporal scale (active regions

during flares or in quiescent conditions, coronal holes, quiet Sun)

Whatever the active region target, the high spatial resolution of HRI is used in a mode close

to A mode with a time resolution of the order of 1 sec. For the polar coronal hole mode, the high

latitude is mandatory and the compression in Lalpha possibly lower than 15 (mode

C). Complementary observations of SPICE in Lbeta would be very useful.

5.1.2 What is the Ly-α emission and absorption in the cool atmosphere (especially in polar

coronal holes)? To what extent is it a function of latitude?

The observing mode is C but the observations must be repeated during the various orbits

(always at perihelion) in order to cover a large range of latitudes. Complementary observations of

SPICE in Lbeta would be very useful.

5.1.3 Study the coronal He abundance

The S mode should be OK since there is overlap between METIS FOV and FSI (304). One

could modify this mode in three ways: not cropping to 4Rs by 4Rs (in order to increase the FOVs

overlap)), decreasing the cadence (down to about an hour) and decreasing the compression by a

factor 5.

5.1.4 Reconstruct the solar EUV irradiance (Haberreiter et al., 2014) for all hemispheric

directions, in particular for higher latitudes

Synoptic F174 (and additionally F304) at a cadence of 1/day; high latitude observations are

of particular interest; coordination with other imagers in space, if possible.

5.1.5 Targets of opportunity (planetary quadratures, comets, …).

For Jupiter and Saturn, the most obvious targets are the aurorae (see observations with the

Hubble Telescope) to be observed in Lalpha. The visibility will depend on the planet-SO distance.

The signal is time variable (from a few kR to MR, Nichols et al. 2007) and detectable (at the typical

distance of 4 AU for Jupiter). The disk would cover about 30*30 pixels of HR Lalpha. This assumes



an out-of-limb offset pointing … To perform further S/N computations, it is necessary to know the

pointing stability.

If one refers to Bennett comet observations (Bertaux, Blamont, Festou 1973) performed at

distances smaller than 1 AU, the Lalpha signal should be about 10^^-6 the solar signal with a very

large field. This requires out-of-limb offset pointing of HRI Lalpha with perhaps a “reconnaissance”

imaging in FSI 30.4 nm. To perform further S/N computations, it is necessary to know the pointing

stability. If SO is offset, METIS could observe the extended bubble of the comet in Lalpha.

5.2 Additional Science Objectives of EPD

5.2.1 Energetic particles in Venus and Earth environment

5.2.2 Galactic and anomalous cosmic rays. Long-term and short-term modulation and spatial

variation.

5.2.3 Energetic neutral atoms (Wang et al., 2014). (How does EPD team mean to address this?)

5.2.4 Jovian electrons

5.3 Additional Science Objectives of MAG

N/A

5.4 Additional Science Objectives of METIS

5.4.1 Hydrogen Ly-α emission by the atmosphere of planets (e.g. Venus, Jupiter) (Chaufray et

al., 2012; Colwell, 1998; Jaffel et al., 2007; Menager et al., 2010).

5.4.2 Study of Sungrazing comets

5.4.2.1 Understand cometary properties and evolution by mapping the hydrogen Ly-α emission,

proportional to the outgassing rate, along its trajectory close to the Sun (Raymond et al., 1998; Uzzo

et al., 2001; Bemporad et al., 2005; Ciaravella et al., 2010).

5.4.2.2 Investigate the fragmentation of the cometary nucleus from the variation with the

heliocentric distance of the outgassing rate and from the splitting of the cometary tail.

5.5 Additional Science Objectives of PHI

(All of the following objectives require simultaneous measurements from another satellite or from

the ground.)

5.5.1 How strongly does the solar luminosity vary and what is the source of these variations?



5.5.1.1 Provide magnetic field and continuum brightness data in order to determine Sun’s irradiance

at different heliospheric latitudes and phases of the solar cycle (Krivova et al., 2003).

5.5.1.2 Comparison with Sun-like stars’ larger irradiance variations (Schatten et al., 1993; Knaack et

al., 2001):

• Is it a line of sight effect relative to their rotation axes or

• Will the Sun also display larger variations in the future?

5.5.2 Other Science Objectives

5.5.2.1 Effect of granulation and oscillations, i.e. interaction of modes and convection (stereoscopic

helioseismology).

5.5.2.2 Two components of velocity. What is the relationship between the components of the

velocities in granulation? Supergranulation? Various modes in quiet Sun? (stereoscopic

helioseismology)

5.5.2.3 Shape of the Sun. Extend oblateness study of (Emilio et al., 2007; Kuhn et al., 2012).

5.6 Additional Science Objectives of RPW

5.6.1 Interplanetary dust: spatial distribution, mass and dynamics.

• Required observations: Electric signatures of dust impact on the spacecraft body (RPW).

5.7 Additional Science Objectives of SoloHI

5.7.1 Interplanetary Dust

• What are the sources and properties of dust in the inner heliosphere?

• Do Sun-Grazing comets contribute to the dust and what is the time dependence?

• Does the scattering function change with heliocentric distance?

• Can the evaporation of dust be detected?

5.7.2 Streamers

• What is the 3-D structure and extent of streamers?

• Can a filamentary nature be detected?

• What is their radial distribution?

• What is the difference between pseudo- and normal streamers?

5.7.3 CMEs

• What is the 3-D structure and extent of CMEs?

• Are flux ropes imbedded in the plasma sheet?

• What are the radial and longitudinal distributions of "blobs", magnetic islands/plasmoids?



5.8 Additional Science Objectives of SPICE

5.8.1 Obtain spectral atlas of representative features in the solar atmosphere.

5.9 Additional Science Objectives of STIX

TBW

5.10 Additional Science Objectives of SWA

5.10.1 What is the temporal variability of 3D thermal particle distributions?

• Burst mode resolution of EAS and PAS will provide fastest thermal distributions recorded at

different radial distances in the solar wind. At high cadence, test:

• how anisotropic and irregular are proton and electron distributions? This can test whether

resonant heating or stochastic heating is the dominant heating mechanism.

• how does the anisotropy change in time / space?

• how are distributions related to local turbulence and global expansion?

5.10.2 How are proton and helium temperatures related to their relative drift, is there any evidence

for resonant heating (Kasper et al., 2013)?

5.10.3 How common are proton beams, where do they come from and what impact do they have on

ambient conditions (wave generation, heating)?

5.10.4 Identify and characterize the various forms of free energy in the particle distribution

functions. Analyze their formation and relaxation processes at sub-second resolution. Characterize

and measure the degree of reversibility and irreversibility in these processes and generalize to the

understanding of the dynamic behavior of weakly collisional media.

5.10.5 Fully characterize the radial and latitudinal variability of 3D thermal particle distributions

under different solar wind conditions (fast, slow wind, in and around CMEs, CIRs, etc.)

5.10.6 Confirm the origin of the electron halo population – is it derived from the scattering of the

strahl (c.f. Maksimovic et al., 2005) and if so what are the physical mechanisms responsible for this

and their radial and latitudinal dependencies?

5.10.7 How do the combined electron strahl and composition signatures vary across reconnecting

solar wind current sheets? Are these consistent with expectations of interchange reconnection

models? What are the combined signatures in other solar wind structures (e.g. shocks, CIRs,

CMEs), and how definitive are these in determining solar origins and/or defining processes

occurring in the solar wind?





4 INSTRUMENT DESCRIPTIONS AND THEIR OPERATIONAL CONSTRAINTS

4.1 EPD

The EPD (Energetic Particle Detector) investigation on Solar Orbiter provides the key in situ

measurements of energetic particles to address the major science objective 3 (How do solar

eruptions produce energetic particle radiation that fills the heliosphere?). EPD also addresses the

other major science objectives by tracking magnetic field lines from the spacecraft to their solar

source, by determining magnetic field line lengths and measuring in situ the interplanetary effects of

evolving Interplanetary Coronal Mass Ejections (ICMEs), and by measuring the latitude and

longitudinal distribution of Solar Energetic Particles (SEPs).

The Energetic Particle Detector suite consists of four sensors measuring electrons, protons, and

heavy ions from helium to iron, and operating at partly overlapping energy ranges from 2 keV up to

450 MeV/n.

The EPD sensors are:

a) Suprathermal Electrons and Protons (STEP)

b) Suprathermal Ion Spectrograph (SIS)

c) Electron Proton Telescope (EPT)

d) High Energy Telescope (HET)

EPT and HET are combined in one unit, and there are two identical EPT-HET units on the

spacecraft.

STEP relies on a magnetic deflection system in one of the two detector units to suppress electrons in

that particular unit. The particles are measured by the use of solid state detectors with ultra-thin

ohmic contacts. SIS applies the time-of-flight by energy technique to measure the composition of

ions. EPT relies on the magnet/foil technique in order to separate electrons from protons and heavier

ions. HET is based on the multiple dE/dx vs. total E technique.

4.1.1 EPD Observables

The relevant SEP characteristics that should be taken under consideration are the following:

1. Energetic Particle Energy

2. Particle Mass

3. Particle Flux Angular information

4. Energetic particle flux temporal information.

These, in terms of instrument parameters, can be written as follows:

a) Energy Range



b) Mass resolution (Ion composition)

c) Fields of View and Angular Resolution

d) Geometric Factor

e) Time Resolution

4.1.1.1 Energy range

The Energy Range is the lowest and highest particle energy that the instrument is capable to

measure. It is defined for the different particle species:

• Electrons (e)

• Protons (p)

• Heavy ions from Helium to Iron

The EPD sensors performance regarding particle energy is given in the following table (Table

4.1.1.1.1) and its energy coverage for the different species in Figure 4.1.1.1.1.



4.1.1.2 Mass resolution

The mass resolution defines both the closest elements, and isotopes, that the instrument is capable to

resolve.

The EPD performance regarding particle mass is given in Table 4.1.1.2.1.

4.1.1.3 Fields of view and angular resolution

The Field of View (FoV) is that portion of the sky that is covered by the cross-sectional area

spanned by the triggering detector and the front defining detector or aperture of the various EPD

sensors. Depending on the EPD sensor, its cross section can be either circular or rectangular.

The angular resolution of a telescope is defined as its ability to distinguish different directions of the

incoming particles.

Here, we will provide information about the number of apertures, both in and out of the ecliptic

plane, the “FoV” sizes and the instrument angular resolution.

The directional information performance of EPD is given in Table 4.1.1.3.1 Figure 4.1.1.3.1 shows

the FoV of EPD referred to the S/C reference frame. Note the clustering of FoVs around the mean

Parker spiral direction.



4.1.1.4 Geometric factor

The geometric factor is defined as the integral over the telescope FoV of the quantity dSd (where

S is the telescope/detector surface and is the solid angle that it subtends). The geometric factor is

the parameter that allows the computation of differential intensities from measured quantities. In

practice, the geometric factor is the factor of proportionality relating the counting rates to the

intensity in the field of view. Geometric factors for each sensor have been optimized after a careful

analysis of the expected fluxes close to the Sun over the very wide energy range covered by EPD.

Table 4.1.1.4.1 shows the geometric factor of the different EPD sensors.



4.1.2 EPD modes

4.1.2.1 EPD Normal mode

TBW

4.1.2.2 EPD Burst mode

TBW



4.2 EUI

The EUI instrument suite is composed of two high resolution imagers (HRI), one at Lyman-α and

one at 174Å, respectively named “HRI_Lyalpha” and “HRI_EUV”, and one dual band full-sun

imager (FSI) working alternatively at the 174 and 304 Å EUV passbands, named “FSI174/304”.

4.2.1 EUI observables

4.2.2 EUI modes and telemetry

TM figures

Allocated TM 20.5 Kbits/s

Download capacity per orbit 6.642 GB = 53.136Gbits

Subtelescopes/units that can be commanded independently

wavelength detector size FOV TM/raw image (15bits) compression rates

FSI 174nm or 304nm 3072 px * 3072 px 3.8ºx3.8º 17.7 MB 6-500

HRI_Lya 121.6nm 2048 px * 2048 px 16.6'x16.6' 7.9 MB 6-50

HRI_EUV 174nm 2048 px * 2048 px 16.6'x16.6' 7.9 MB 6-50

Low Latency programs

Observing mode Instrument Rebin/subfield #Images/h Data Rate Data Rate Max time

(Gbits/h) (Kbits/s) (h)

EUI/FSI Beacon mode (B) FSI 174 768x768 4

0,0005 0,146 unlimited

FSI 304 768x768 4

EUI/HRI Beacon modes (HB) HRI_EUV 2048x2048 0,0417

0,0001 0,029 unlimited

HRI_Lyalpha 2048x2048 0,0417

(1 set/day)

Note that the Beacon mode will run in parallel to most of the FSI programs defined below.

Main science programs

Observing mode Instrument Rebin/subfield #Images/h Data

Rate

Data

Rate

Max time

***


FSI Synoptic mode (S) *

FSI 174 1536x1536 4

0,0056 1,56

unlimited

FSI 304 1536x1536 4



FSI Reference Synoptic mode (R)

* FSI 174 3072x3072 0,0417

0,00022 0,061 unlimited

FSI 304 3072x3072 0,0417

(1 set/day)

FSI Global eruptive event mode

(G) **

FSI 174 or

304 3072x3072 360 3,8 1054,7 12 h

FSI Faint High Corona (FHC) ** FSI 174 or

304 3072x3072 4 0,021 5,86 unlimited

EUV & LYA Coronal hole modes

(C)

HRI_EUV 2048x2048 120 1,5 416,7

30 h

HRI_Lyalpha 2048x2048 120

EUV & LYA Quiet Sun modes (Q) HRI_EUV 2048x2048 450 14,1 3906,3

3 h


EUV & LYA Active Region modes

(A)

HRI_EUV 2048x2048 1800 16,9 4687,5

3 h


* Note that all FSI modes will typically be combined with the Beacon mode (B) defined above to

assure regular EUI LL data. The data rates above do not yet include this extra TM, total data rate is

given in the mode-specific pages (see e.g. FSI Synoptic mode (S)).

** Some more specific FSI modes will have both Beacon mode (B) and Synoptic mode (S) running

in parallel. The data rates above do not yet include this extra TM, total data rate is given in the

mode-specific pages (see e.g. FSI Global eruptive event mode (G)).

*** Maximum time for this mode to run per orbit is based on the average EID-A rate over the 30

days RS windows. There will be times though that more downlink is available, and at other times

more stringent restrictions will apply.

Additional science programs

Observing mode Instrument Rebin/subfield #Images/h Data Rate Data Rate Max time


EUV & LYA Eruptive Event modes

(E)

HRI_EUV 2048x2048 3600 22,5 6250

2 h


EUV & LYA Discovery modes (D) HRI_EUV 2048x2048 3600 22,4 6225

2 h




In the following sections, we present in detail the EUI modes.

4.2.2.1 FSI Beacon mode (B)

This mode will be used with FSI only to produce (part of) the EUI Low Latency data. It generates

highly compressed synoptic data, with extremely restricted TM.

• Parameters

TBC

• Resource usage

Observing

mode Target Instrument Channel Cadence Rebinning Compression

Calculated

TM #Images/h Data Rate

Data

Rate

Max

time

(nm) (s) or subfield rate (bits/image) (Gbits/h) (bits/s) per orbit

FSI

Beacon

mode (B)

= highly

compressed FSI 17,4 900 4 125 70778.88 4 0.00053 157.3 unlimited

synoptic

data FSI 30,4 900 4 125 70778.88 4

• Maximum time in this mode

unlimited

whole orbit would take ~4% of allocation

30 days takes 0.8%

4.2.2.2 FSI Synoptic mode (S)

This mode is using the FSI full disk telescope and provides synoptic full disk images (full sun up to

4x4 solar radii).

Depending on the distance from the sun, images are either rebinned 2x2 (close to Sun) or a quarter

of the detector is read out (subfield, far from Sun). Both choices result in same data rate.

In practice, this EUI mode will run in parallel to FSI Beacon mode (B) to ensure availability of

regular EUI/FSI Low Latency data.

• Parameters

Cadence and compression rate (TBC)



• Resource usage

Observing mode Target Instrument Channel Cadence Rebinning Compression

rate #Images/h

Data

Rate

Data

Rate

Max time per

orbit

(nm) (s) or

subfield (Gbits/h) (bits/s)

Synoptic mode (S) Full sun up

to 4x4R

FSI 17,4 900 2x2 46.88 4 0.0056 1677.7

unlimited

FSI 30,4 900 2x2 46.88 4

FSI Synoptic + FSI

Beacon (S+B) 0.0062 1835.0 unlimited


unlimited

3 RSWs = 10% allocation

full orbit (168 days) = 56% allocation

4.2.2.3 FSI Reference Synoptic mode (R)

This mode is using the FSI full disk telescope and provides synoptic full disk images (full sun up to

4x4 solar radii).

Full sun images are obtained with better resolution and less compressed than in FSI Synoptic mode

(S). Typically, only 1 set of those is planned per day (cadence 1/24hrs).

• Parameters

Cadence and compression rate (TBC!!)

• Resource usage

Observing


Calculated


Data

Rate Max time

(nm) (s) or subfield rate (bits/image) (Gbits/h) (bits/s) per orbit

FSI

Reference

synoptic (R)

4x4Rsun

window

FSI 17,4 86400 2x2 12.5 30670848 0,0417

0.00022 65.5

unlimited

FSI 30,4 86400 2x2 12.5 30670848 0,0417

FSI

Reference

Synoptic

(R) + FSI

Beacon (B)

0.00075 222.8 unlimited



• Maximum time in this mode (R)

unlimited

30 image duos (1 per each RSW day) = 0.01% allocation

full orbit (168 days) = 0.07% allocation

4.2.2.4 FSI Global eruptive event mode (G)

This mode is using the FSI full disk telescope and aims at capturing a global event.

• Parameters

Channel, cadence and compression rate (TBC!!)

• Resource usage

Observing


Calculated


Data

Rate

Max

time

(nm) (s) or subfield rate (bits/image) (Gbits/h) (bits/s) per

orbit

Global eruptive

event 174

(G174)

Full

event FSI 17.4 10 1 10 11324621 360 3.80 1132462 12 h

Global

eruptive event 304

(G304)

Full event

FSI 30.4 10 1 10 11324621 360 3.80 1132462 12 h

Global

eruptive

event 174

(G174)

+ FSI

Synoptic

(S) + FSI

Beacon (B)

3.803 1134140 12 h

Global

eruptive

event 304

(G304)

+ FSI

Synoptic

(S) + FSI

Beacon

(B)

3.803 1134140 12 h


restricted to 12 hours per orbit = 92% of allocation



4.2.2.5 FSI Find Event mode (FE)

This mode is using the FSI full disk telescope to detect events onboard. The FSI thumbnails used for

event detection will not be brought to ground so the only TM resource usage in this mode is the

generation of parallel synoptic imagery and beacon (LL) imagery.

• Parameters

TBW

• Resource usage

Observing mode Target Instrument Channel Cadence Rebinning Compression #Images/h Data

Rate

Data

Rate

Max

time

(nm) (s) or

subfield rate (Gbits/h) (bits/s) per orbit

FSI Find Event (FE174) Event

detection FSI 0 0 unlimited

FSI Find Event (FE174)

+ FSI Synoptic (S) + FSI

Beacon (B)

0.006 1835 unlimited


unlimited

4.2.2.6 FSI Faint High Corona mode (FHC)

This mode is using the FSI full disk telescope, with the disk occulter in the FSI filter wheel. It aims

at capturing the faint high corona off-limb.

• Parameters

Channel, cadence and compression rate (TBC)

• Resource usage

Observing mode Target Instrument Channel Cadence Rebinning Compression Calculated

TM #Images/h

Data

Rate

Data

Rate

Max

time

(nm) (s) or

subfield rate (bits/image) (Gbits/h) (bits/s) per orbit

Faint High Corona

174 (FHC174)

Full

event FSI 17.4 3600 1 6.25 22649241.6 1 0.0211 6291.5 unlimited

Faint High Corona

304 (FHC304)

Full

event FSI 30.4 3600 1 6.25 22649241.6 1 0.0211 6291.5 unlimited






Faint High

Corona 174

(FHC174)

+ FSI Synoptic

(S) + FSI Beacon

(B)

0.0267 7969.0 unlimited

Faint High

Corona 304

(FHC304)

+ FSI Synoptic

(S) + FSI

Beacon (B)

0.0267 7969.0 unlimited


unlimited

the combined mode FHC+S+B, running through all RS windows (30 days) would consume ~40% of

the orbital allocation

4.2.2.7 EUV & LYA Beacon modes (HB)

This is HRI's Low Latency programme that produces one set of small FOV images per day.

The mode can be commanded on HRI_EUV and HRI_Lya telescopes independently.

• Parameters

TBW

• Resource usage

Observing

mode Target Instrument Channel Cadence Rebinning Compression #Images/h

Data

Rate

Data

Rate

Max

time

(nm) (s) or

subfield rate (Gbits/h) (bits/s) per orbit

EUV Beacon

mode

Can be run on any target

HRI_EUV 17,4 86400 1 47 1/day 0.00005 15.5 unlimited

LYA Beacon

mode HRI_Lyalpha Ly alpha 86400 1 47 1/day


unlimited

4.2.2.8 EUV & LYA Coronal hole modes (C)

This mode uses HRI telescopes targeting a full coronal hole, a CH boundary or polar plumes.


• Parameters



cadence, compression TBC

• Resource usage

Observing


Calculated

TM #Images/h

Data

Rate

Data

Rate

Max

time

(nm) (s) or

subfield rate (bits/image) (Gbits/h) (bits/s)

per

orbit

EUV Coronal

hole (C)

Full CH +

boundary HRI_EUV 17,4 30 1 6.25 10066329.6 120 1,5 447392.4 30 h

LYA Coronal

hole (C) + plumes HRI_Lyalpha Ly alpha 30 1 18.75 3355443.2 120


restricted to 30 hours per orbit = 90% of allocation (if both telescopes are used), if all resulting TM

gets downlinked

4.2.2.9 EUV & LYA Quiet Sun modes (Q)

This mode uses HRI telescopes targeting a region of quiet sun. It is most appropriate to be used

close to perihelion, near co-rotation.


• Parameters

TBW

• Resource usage

Observing

mode Target Instrument Channel Cadence Rebinning Compression #Images/h

Data

Rate

Data

Rate

Max

time

(nm) (s) or

subfield rate (Gbits/h) (bits/s)

per

orbit

EUV Quiet Sun (Q)

Full HRI FOV

of quiet sun

HRI_EUV 17,4 8 1 9.4 450

14.06 4194304.0 3 h LYA Quiet Sun

(Q) HRI_Lyalpha Ly alpha 1 1 18.75 3600


TM-hungry mode, limited to 3 hours per orbit = 85% of EID-A allocation (if both telescopes are

used), if all resulting TM gets downlinked!



4.2.2.10 EUV & LYA Active Region modes (A)

This mode uses HRI telescopes targeting an active region on the sun. It is most appropriate to be

used close to perihelion, near co-rotation.


• Parameters

TBW

• Resource usage


Rate

Data

Rate

Max

time

(nm) (s) or


per

orbit

EUV Active

Region (A) Full HRI FOV of active region

HRI_EUV 17,4 2 1 18.75 1800

16.88 5033165 3 h LYA Active

Region (A) HRI_Lyalpha Ly alpha 1 1 18.75 3600


TM-hungry mode, limited to 3 hours per orbit (if both telescopes are used), if all resulting TM

gets downlinked!

4.2.2.11 EUV & LYA Eruptive Event modes (E)

This mode uses both HRI telescopes targeting an eruptive event on the sun. It is most appropriate to

be used close to perihelion, near co-rotation.


• Parameters

• Resource usage


Rate

Data

Rate

Max

time

(nm) (s) or


per

orbit

EUV Eruptive event

(E)

Full FOV on event

HRI_EUV 17,4 1 1 18.75 3600

22.5

6710886

2 h LYA Eruptive event

(E) HRI_Lyalpha Ly alpha 1 1 18.75 3600




TM-hungry mode, limited to 2 hours per orbit = 98.5% of allocation, if all resulting TM gets

downlinked!

4.2.2.12 EUV & LYA Discovery modes (D)

This mode uses HRI telescopes, with the Ly Alpha channel in its highest cadence 0.1s, to capture

high cadence dynamics outside of flaring times. HRI EUV is used with full FOV, in Lyalpha only

a subfield is read out. This mode is most appropriate to be used close to perihelion, near co-rotation.


• Parameters

• Resource usage


Rate

Data

Rate

Max

time

(nm) (s) or


per

orbit

EUV Discovery

(D)

High cadence

dynamics

HRI_EUV 17,4 1 1 18.75 3600

22.5 6683643 2 h LYA Discovery (D)

HRI_Lyalpha Ly alpha 0,1 3,175 18.75 36000


TM-hungry mode, limited to 2 hours per orbit = 90% of allocation (for both telescopes together), if

all resulting TM gets downlinked!



4.3 MAG

The measurement of magnetic fields in space science is constrained by environmental

factors such as temperature and radiation, and also by restrictions on resources available for the

instrument such as weight and power. The fluxgate magnetometer principle has emerged as an

optimum compromise method for magnetic field measurement in space, since such instruments are

rugged, low in power and mass and offer high precision.

The basic operating principle of the fluxgate magnetometer is well known and documented.

A soft magnetic core, usually toroidal in shape, is wound with a coil and driven into saturation with

an AC excitation current. The external magnetic field (which is to be measured) distorts the

symmetry of the magnetic flux in the core, which generates a signal at even harmonics of the

excitation frequency. This signal is detected by a sense coil wound around the core. By feeding back

a current into the sense coil proportional to the measured signal, the ambient field is backed-off and

the sensor operates in null-mode, thereby improving linearity.

Typically, the second harmonic of the drive frequency is notch-filtered, amplified,

synchronously detected, integrated, and used to drive the feedback current. It is this current which is

proportional to the ambient magnetic field. The output of a traditional analogue magnetometer is the

voltage used to drive this feedback current and is usually digitized by an analogue-to-digital

converter. In such a design, the signal processing to extract the second-harmonic signal is performed

by an analogue circuit. This is an efficient and well-developed method, however the filtering and

synchronous detection stages all show temperature dependencies, and the analogue component

count and mass are not negligible. By replacing the analogue processing with a digital system, this

signal processing is performed as a software function in the digital domain. As well as eliminating

some of the temperature effects, and reducing mass, this provides an inherent flexibility in the

design, since the processing parameters can be modified by software.

The MAG instrument performance is summarized in the following table:



4.3.1 MAG observables

4.3.2 MAG modes

4.3.2.1 MAG Normal mode

Burst modes run in parallel with normal modes. Here, HK and LL TM are included in the normal

mode data rate.

Mode IBS Cadence (vps) OBS Cadence (vps) Data Rate (bps) Note

Normal 1 16 1000 Default Option

Equal8 8 8 950

Low 1 1 250

4.3.2.2 MAG burst mode

Burst modes run in parallel with normal modes. Here, HK and LL TM are included in the normal

mode data rate. Burst mode rates here represent the additional TM from burst data. In principle, any

normal mode can be combined with any burst mode, although not all combinations make scientific

sense unless selective downlink is implemented for MAG.

Mode IBS Cadence (vps) OBS Cadence (vps) Data Rate (bps) Note

Burst 8 128 6892 Normal + Burst = 7892

Burst1 1 128 6542 Normal + Burst1 = 7542

Equal16 16 16 1692 Normal + Equal16 = 2692

Equal128 128 128 12924 Normal + Equal128 = 13924



4.4 METIS

METIS is an inverted-occultation coronagraph for the simultaneous observation of the whole

Sun corona in the visible light and around the UV wavelength of the H I Lyman- spectral line.

METIS will be capable of obtaining for the first time simultaneous imaging of the full corona in

polarized visible light (580-640 nm) and narrow-band ultraviolet H I Ly-α (121.6±10 nm). These

measurements will allow a complete characterization of the most important plasma components of

the corona and the solar wind, i.e. electrons and protons.

METIS can simultaneously image the visible and ultraviolet emission of the solar corona

and diagnose, with unprecedented temporal coverage and spatial resolution (down to about 4000

km), the structure and dynamics of the full corona in the range from 1.6 to 3.0 solar radii (R☉) at

minimum perihelion (0.28 AU), and from 2.8 to 5.5 R☉ at 0.5 AU. This region is crucial in linking

the solar atmosphere phenomena to their evolution in the inner heliosphere, and the study of its

properties is very important in meeting the Solar Orbiter fundamental science goals.

The instrument has been conceived to perform off-limb and near-Sun coronagraphy, with the

aim of addressing the three key scientific issues concerning:

• the origin and acceleration of the fast and slow solar wind streams;

• the origin, acceleration and transport of the solar energetic particles;

• the transient ejection of coronal mass (coronal mass ejections, CMEs) and its evolution in the

inner

heliosphere.

In addition, METIS can contribute to the study of the properties and evolution of Sun-

grazing comets. It is also able to characterize the properties of the coronal regions crossed by the

Solar Probe Plus spacecraft, during

its transit close to the Sun.

4.4.1 METIS observables

The expected science products of the off-limb and near-Sun coronagraphy performed by METIS

are:

• global maps of VL coronal emission in the range 580-640 nm for each of the 4 polarization states

of the

Liquid Crystal Variable Retarder

• global maps of the H I Ly-α (121.6 nm) UV coronal emission.

The above measurements can be obtained simultaneously on board.

On ground, the VL channel measurements can be combined to obtain maps of either the total and/or

the

polarized brightness (tB/pB) of the corona.

In addition, 8 VL light curves, 1 for each sector of the field of view, can be produced (at the

moment, only as

Low Latency data during Coronal Mass Ejections observations).



Further processing of calibrated Ly-α and pB images can provide:

• global maps of electron densities, from the pB maps;

• global maps of the outflow velocity of the H component of the solar wind (by Doppler dimming

technique.

4.4.2 METIS modes

Note that all data volumes & rates are based on powers of 10, i.e. 1Gbit = 109 bits.

TM figures


Download capacity per orbit 3.4 GB = 27.2 Gbits


wavelength detector size FOV TM/raw image compression rates

METIS_VL polarized VL in 580-640nm (4 images) 2048 px * 2048 px annular FOV 1.5º-2.9º 7.34 MB (14bit/px) 13-50

METIS_UV 121.6nm (Ly_beta) 1024 px * 1024 px annular FOV 1.5º-2.9º 1.84 MB (14bit/px) 13-50

Full set of raw METIS images (4 VLs+1UV) accounts for 31.20MB.

If partially* rebinned and compressed (factor 2.5), a full METIS dataset accounts for 4.6MB to

16MB depending on the science goal.

* METIS images are rebinned differently in the inner band of the FOV as in the outer band:

typically, with binning factor 2 < 2.5º and 4 >2.5º for global corona and binning factors 2 < 2.5º

and 4 >2.5º for smaller scale features like CMEs and oscillations

Observational modes


Observing mode Use case Telescope(s) Cadence Data Vol /

set

Data

Rate

Max

Duration

(min) (Mbits) (Kbits/s)

METIS standard

modes

WIND

electron density measurement and solar wind

velocity

<0.5AU (1.6 - 5 solar radii)

METIS_VL+UV 5-30min

14.12Mb

(4VL+1UV)

8 - 47 days

MAGTOP

slow wind source regions (magnetic topology)

near perihelion (1.6 - 5 solar radii)

METIS_VL+UV 5-20min 23.34Mb

(4VL+1UV) 20 - 78 days



GLOBAL

Global corona configuration/evolution during

CME events

close or far from sun, binning dependent on

distance

METIS_VL+UV 5-30min 12Mb - 8Mb

(4VL+1UV) 4.5 - 40

days

LT-CONFIG

Long-term evolution of coronal configuration

close or far from sun, binning dependent on

distance


(4VL+1UV) 2.8 - 4.2 days

METIS special

modes

FLUCTS+TBF

Brightness fluctuations spectra

at perihelion only (1.6 to 3 solar radii)

METIS_VL only

3 step

acquisition,

variable cadence:

1s - 20s -

10min

4.67Mb

(1 VL image)

311

(avg/hr)

hours

CMEOBS

CME driven shocks and SEP (filamentary

structure)

triggered by METIS-internal flag (at any

distance)


(4VL+1UV) 77.8 - 389 hours

COMET

Mapping emission of sungrazing comets

(at any distance)


(4VL+1UV) 20 - 78 days

PROBE

Coordinated observations with Solar Probe Plus

(at any distance)


(4VL+1UV) 11 - 66 days

METIS observing mode characteristics

In the following definitions of METIS observing modes, the following naming conventions are

applied:

1. Size of raw frames (sets of samples representing the digitized signal given by each pixel of

the relevant detector):

1. VLD: (2048x2048) @ 14 bits (analogue mode; spatial scale per pixel: 10”);

2. UVD: (1024x1024) @ 14 bits (analogue mode; spatial scale per pixel: 20”).

The UVD can be operated in photon counting mode also; the photon counting mode

is mainly required to overcome possible degradation of the UV detector during a



long-lived mission such as Solar Orbiter and to ensure the capability of observing

also in such a case.

2. DIT: Detector Integration Time, i.e., the time interval among which all detector pixels are

collecting photons. It represents the actual exposure of a single read out of the detector

(frame).

3. NDIT: number of detector integrations, i.e., the number of frames to be averaged in order to

obtain an “Acquisition”; this is the image obtained as result of averaging, pixel by pixel,

NDIT frames.

4. TACQ: Acquisition Time, the overall integration time corresponding to a single acquisition.

The values are set in order to get a limited number of events (cosmic rays) to be removed.

5. Cosmic ray (CR) removal: software procedure to clean the acquired images from spurious

signal given by the cosmic rays and SEP.

6. CME flag: software procedure to automatically trigger the rising of a CME event.

7. NACQ: number of acquisitions to be averaged in order to obtain an “Exposure”, i.e., the

image obtained as result of averaging, pixel by pixel, NACQ acquisitions.

8. TEXP: Exposure Time, i.e., the overall time corresponding to a single Exposure.

Single acquisitions will be averaged over the exposure time interval, unless a CME flag

occurs. Exposure times are determined considering the count rates estimated for the Sun at

its minimum of activity (Tables 6-8), as expected for a launch date of July 2017. The values

are set:

1. in order to get the best performance of the detector for the brightest features in the

expected image and, contemporarily,

2. get still a significant signal-to-noise ratio for the faintest features in the same

expected image;

3. on the basis of the typical life-times of the coronal structures considered;

4. taking into account the need of limiting the data volume of the final science

images downloaded to the ground.

9. NPOL: number of polarization (LCVR) angles used (cycled) during the scientific

measurement.

10. TOH: Overhead Time, given by the instrument in order to perform the scientific measurement

(i.e., synchronization, LCVR commanding, etc.).

11. TW: Waiting Time, defined by the operator within the observation timeline.

12. TCAD: Cadence Time, needed to get NPOL VL images or 1 UV image (within the same

scientific measurement, the time between two consecutive Exposures).

13. The counts per pixel in the final science images (to be downloaded by TM) are average

values calculated over the summed-up acquisitions and the pixels binned together. In this

case the depth of each final science image is 14 bits.

The acquisition time relationships are

VLD UVD



TACQ = NDIT * DIT = DIT , NDIT = 1 TACQ = NDIT * DIT

TEXP = NACQ * TACQ TEXP = NACQ * TACQ

TCAD = NPOL * TEXP + TOH + TW TCAD = TEXP + TOH + TW

The characteristic requirements of each observing mode are given in the two following

sections METIS standard modes and METIS special modes.

METIS standard modes

source: Metis User Manual Iss3 (May 2016) - table 4.1.3.2

Specific instrument observing modes have been defined in order to address the scientific objectives

of the METIS investigation. In general, METIS observations consist of global maps of the coronal

emission in UV H I Ly-α and VL (580-640 nm range), obtained with different spatial resolution and

detector exposure time, depending on the science goal and the instantaneous field of view (FoV).

More in detail, the observing modes defined for METIS are:

WIND – Measurement of the electron density and the solar wind outward expansion velocity

Fast and slow solar wind streams are identified in the global maps according to the values of the

outflow velocity of the H component.

MAGTOP – Wind outflow velocity measurements and relationship with the magnetic topology

Maps of the outflow velocity of the H component along streamer/coronal hole interfaces, above

streamers cusps and inside streamers. Relationship with the coronal magnetic configuration.

GLOBAL – Global corona configuration/evolution measurements before, during and after CME

events These measurements provide the geometry of the neutral H and e- corona and its evolution in

time, giving information on the timing, mass content and overall dynamics of coronal mass

ejections. They are also crucial to measure the directionality of the plasma erupted from the Sun, in

order to infer its geo- effectiveness and predict the impact on the near-Earth environment. The

evolution related to the CME transit can be followed out to the orbit of the Solar Probe Plus

spacecraft (~9 R! at 0.8 AU).

LT-CONFIG – Long-term evolution of the coronal configuration

These measurements are used to monitor the evolution of the large-scale corona and, during out-of-

ecliptic observing windows, to determine the longitudinal distribution and evolution of the electron

density in the solar corona, as well as of the mass and energy flux carried away by the solar wind.



METIS special modes

source: Science Performance Document Iss 4.0 (CDR version - March 2016) & Metis User Manual

Iss3 (May 2016) - table 4.1.3.2

Specific instrument observing modes have been defined in order to address the scientific objectives

of the METIS investigation. In general, METIS observations consist of global maps of the coronal

emission in UV H I Ly-α and VL (580-640 nm range), obtained with different spatial resolution and

detector exposure time, depending on the science goal and the instantaneous field of view (FoV).

The special observing modes defined for METIS are:



FLUCTS and TBF – Brightness fluctuations spectra (best near perihelion) - acquisition sequence

details below

High spatial resolution and cadence time series of VL coronal brightness over a range of distances

covering from 1.6 R! out to 5 R!, to constrain the amplitude of the density fluctuations spectrum.

High frequency: TCAD=1 s (FLUCTS), =20 s (TBF).

The typical duration of this program is approximately 1 hour

CMEOBS – CME propagation, related driven shocks and filamentary structures, SEP accelerated

by CMEs Measurement of the electron density and outward expansion velocity gradient at high

spatial resolution and temporal cadence. This mode is activated by a proper CME event flag.

Such measurements can be used to identify the path of the shock front where particles can be

accelerated in the outer solar corona. Moreover, combined with radio observations, they can help to

distinguish flare- accelerated SEPs from those associated with CMEs.

COMET – Mapping the emission of Sungrazing comets

These measurements are used to monitor the evolution of the cometary emission along its trajectory

close to the Sun.

PROBE – Coordinated observations with Solar Probe Plus (SPP)

Characterizing the properties of the coronal regions crossed by the SPP spacecraft, during its transit

close to the Sun.



FLUCTS + TBF mode: Coronal brightness fluctuations at high frequency: acquisition details

Source: Science Performance Report 4.0

The program relevant to the study of the high frequency brightness fluctuations spectra (observing

modes FLUCTS and TBF, see Table 12) will be carried out in about 1 hr, as shown in Fig. 2:

The program relevant to the study of the high frequency brightness fluctuations spectra (observing

modes FLUCTS and TBF, see Table 12) will be carried out in about 1 hr, as shown in Fig. 2:

- Run mode FLUCTS twice as follows:

• step 1: acquisition of ~60 consecutive raw VL images with DIT=1 s, NPOL=1, temporarily

stored in the instrument memory

• step 2: wait for on-board processing (masking, binning, compression, see Sect. 5.1.1) of the

raw images (TW ~10 s per image, in total ~600 s) and transfer of the reduced size images to

the S/C SSMM (280.2 Mb) to empty the instrument memory



- Change from mode FLUCTS to mode TBF

- Run mode TBF:

• acquisition of ~120 VL images with DIT=20 s, NPOL=2, changing the polarization angle by

90° after 10 s during each acquisition; on-board processing of image i and transfer to the S/C

SSMM is performed contemporarily to the acquisition of image i+1 (data volume: 560.4

Mb)

The expected data volume produced by the instrument in 1 hr is equal to 1120.8 Mb.



4.5 PHI

PHI is based on two telescopes, a High Resolution Telescope (HRT) and a Full Disk Telescope

(FDT) which define the targeted FOV. The polarimetric modulation is carried out by two

polarization modulation packages PMP, each one based on two temperature stabilized liquid crystal

variable retarders (LCVRs) located close to an intermediate

focus of each telescope. In order to allow for high-resolution polarimetric observations the HRT is

equipped with

an Image Stabilization System (ISS) based on a closed loop correlation tracker. This system is

designed to reduce the image motion on the science detector by moving a tip/tilt mirror. The

correction signals are obtained from cross-correlating images of the observation target obtained with

a fast active pixel sensor (APS) camera. The two telescopes sequentially – adjustable with a Feed

Select Mechanism (FSM) – feed the FG which is based on a tunable solid state LiNbO3 Fabry-Perot

etalon and an order sorting prefilter (OSPF). Tuning the transmission bands of the etalon is carried

out by changing the refractive index of the LiNbO3 wafer by means of adjusting the applied high

voltage which is produced by a High Voltage Power Supply (HVPS). Finally, the solar light will be

focused on two science detectors, each one based on 2kx2k APS sensors. In order to reduce the

power and radiation entry into the instrument both telescopes are equipped with Heat Rejection

Entrance Windows (HREWs) mounted at the heat shield of the S/C. For calibration purposes both

the HRT and the FDT as well as the ISS are equipped with re-focusing mechanisms (RFMs).

Instrument control, data pre-processing (calibration and polarimetric demodulation) as well as

radiative transfer equation (RTE) inversion will be carried out in a Digital Processing Unit (DPU)

which is based of two re-configurable and one fixed field programmable gate arrays FPGA. The

DPU is equipped with a large mass memory (4TBytes) which allows for storing calibration data and

science data which cannot be processed in real time.

4.5.1 PHI observables

PHI is an imaging spectro-polarimeter in order to measure the photospheric continuum intensity, Ic,

the full magnetic field vector, B = [|B|, γ, φ] and the line of sight (LOS) flow velocity, v_LOS, in

each point of dedicated field of view (FOV) on the solar disk.

4.5.2 PHI modes

TM figures



Subtelescopes

PHI has 2 telescopes but these use the same detector so at each moment in time only 1 of the

telescopes can observe and generate data!



wavelength detector size FOV Vol/raw

dataset

TM/processed

image

FDT 617.3nm line (measured at 5 wavelengths +

continuum)

2048 px * 2048

px 2ºx2º 302 MB 12 MB

HRT 617.3nm line (measured at 5 wavelengths +

continuum)

2048 px * 2048

px 16.8'x16.8' 302 MB 12 MB

Observational modes

For PHI, there is only 1 data acquisition mode for each telescope. However, the data processing

onboard differs in the following 'observing modes':

Out of date!

To be updated! PHI observational modes have changed recently

Observing mode (=processing

options) Use Case Telescope Cadence Data Rate Max time / orbit

(s) (Kbits/s) (restricted by TM

alloc.)

PHI science mode 0

nominal mode FDT/HRT 60 1607,7824 9 h

e.g. magnetoconvection

PHI science mode 1 local helioseismology FDT/HRT 60 21,81 28 days

Low resol, high cadence target: 100days

PHI science mode 2 magnetic field evolution FDT/HRT 300-3600 3,79-

52,43 ~24 days - orbit

med resol, med cadence

PHI science mode 3 subfield FDT/HRT 60-3600

(network) 1-262 2 days - unlimited

high resol, high/med cadence 300-3600 (AR)

PHI science mode 4 photospheric context FDT/HRT 300 182-210 2-3 days

PHI science mode 5

global helioseismic/synoptic observations

FDT 60-3600 1,5-5,5 100 days or more

PHI science mode 6 daily context FDT only 1-4 per day 0,24-1 unlimited

( SYNOPTIC MODE)

PHI science mode 7 burst/flare mode, triggered by STIX HRT 60 350 ~40 hours



Science modes specific for Helioseismology

TBW

Low Latency or Precursor programs

Observing mode Use Case Telescope Cadence Data Volume Max time / orbit

PHI Full-disk Precursors Full sun continuum + magnetogram FDT or HRT 1set/day

3,2 Mibits/set

= 0,2 MB / set

PHI Calibration set full calibration set FDT and HRT n/a

8,63 Mibits/set

= 1,1 MB / set

Special datasets

Observing mode Use Case Telescope Cadence Data Volume Max time / orbit

Raw data downlink full, raw image set FDT or HRT n/a

2304Mibits / set

= 302 MB / set

1 set = 5% orbital allocation

PHI Full-disk Precursors

PHI can take full disk precursor images in the days before a RS window. This 'mode' is configured

to generate those, once a day.

• Parameters

TBD

• Resource usage

Data

processi

ng

modes

Use case

Dataset

after

inversi

on

Rebinni

ng

digital depth/phys. quantity

(bit/px)

Min

Caden

ce

Max

Caden

ce

Compressi

on

Data

Vol/set

Min

TM

Rate

Max

TM

Rate

Rate Rate

HRT/FD

T Mibits

or

subfield

I_

c

B_LO

S

v_LO

S

gam

ma

ph

i (s) (s) Rate (*) Mibits

(kibits/

s)

(kibits/

s)

(kibits/

s)

(kbp

s)

Full-disk

Precurso

rs

full sun

continuum 640,00 2 8 8 0 0 0 86400 86400 10 1,60 0,037 0,037 0,04

+magnetogr

am 640,00 2 8 8 0 0 0 86400 86400 10 1,60



FDT or HRT

1set/day

=0.2MB/set

PHI science mode 0

PHI science mode 0 is the nominal mode of PHI. However, it generates that much TM that it cannot

all be downloaded. Only 9 hours of Mode 0 fit in PHI's orbital allocation.

This mode can be run either with the Full Disk Telescope (FDT) or with the High Resolution

Telescope (HRT), each projecting onto the same camera plane.

• Parameters

TBC if any

• Resource usage

Data processing

modes

Use case Dataset after

inversion

Rebinning digital depth/phys. quantity (bit/px)

Min Cadence

Max Cadence

Compression Data Vol/set

Min TM Rate

Max TM

Rate

Rate Rate

HRT/FDT Mibits or

subfield I_c B_LOS v_LOS gamma phi (s) (s) Rate (*) Mibits (kibits/s) (kibits/s) (kibits/s) (kbps)

PHI

science

mode 0

nominal mode 640,00 1 10 10 10 8 8 60 60 2 92,00 1570,1 1570,1 1570,1 1607,7824

e.g. magnetoconvection

PHI science mode 1

PHI science mode 1 is used for local helioseismology. It has lower resolution than PHI science

mode 0 (rebinned or sub fielded to 512x512) but still high cadence of 1 processed dataset per

minute (consists of 5 physical value images processed from 24 raw).

PHI mode 1 is designed to run both on HRT and FDT.

• Parameters

TBC if any

• Resource usage

Data

processing

modes

Use case

Dataset

after

inversion

Rebinning digital depth/phys. quantity

(bit/px)

Min

Cadence

Max

Cadence Compression

Data

Vol/set

Min TM

Rate

Max TM

Rate Rate Rate

HRT/FDT Mibits or


PHI science

mode 1

local

helioseismology 640,00 4 0 0 10 0 0 60 60 2 1,25 21,3 21,3 21,3 21,81

Low resol, high cadence

(512x512)



PHI science mode 2

PHI science mode 2 has lower resolution than PHI science mode 0 (rebinned or sub-fielded to

1024x1024) and a medium cadence that can vary from 1 processed dataset per 5 minutes to 1

dataset per hour.

It can be used with both FDT and HRT telescopes.

• Parameters

Cadence (TBC)

• Resource usage

Data

processing

modes

Use case

Dataset

after

inversion


(bit/px)

Min

Cadence

Max

Cadence Compression

Data

Vol/set

Min TM

Rate

Max

TM

Rate

Rate Rate

HRT/FDT Mibits or subfield I_c B_LOS v_LOS gamma phi (s) (s) Rate (*) Mibits (kibits/s) (kibits/s) (kibits/s) (kbps)

PHI

science

mode 2

magnetic

field

evolution

640,00 2 10 10 10 0 0 3600 300 2 15,00 4,3 51,2 1->51.2 52,43

(max)

med

resol,

med

cadence

640,00 2 0 10 0 8 8 3600 300 2 13,00 3,7 44,4 3,79

(min)

(1024x1024)

PHI science mode 3

PHI science mode 3 is a subfield mode with high resolution (full resolution in subfield of

1024x1024 or 512x512) and a medium to high cadence that can vary from 1 processed dataset per

minute (for network studies and ARs) to 1 dataset per hour.

It can be used with both FDT and HRT telescopes.

• Parameters

Subfield size, cadence, physical quantities (TBC)

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOPSD/PHI+science+mode+0



• Resource usage

Data processing

modes

Use case Dataset after

inversion


(bit/px)

Min

Cadence

Max

Cadence Compression

Data

Vol/set

Min TM

Rate

Max TM

Rate Rate Rate

HRT/FDT Mibits or


network/AR

PHI science

mode 3 subfield 640,00 2 10 10 10 0 0 3600 60 2 15,00 4,3 256,0 1->256

262,14

(max)

high resol, high/med

cadence

640,00 4 10 10 10 0 0 3600 300 2 3,75 1,1 12,8 1,09 (min)

(subfield) AR

extra

possibility: 640,00 0 10 0 8 8

PHI science mode 4

PHI science mode 4 is used for photospheric context. It outputs a subset of physical quantities (3

output images out of 5) at a cadence of 5 minutes and maximal resolution.

It can be used both with HRT and FDT.

• Parameters

physical quantities (TBC)

• Resource usage

Data

processing

modes

Use case

Dataset

after

inversion


(bit/px)

Min

Cadence

Max

Cadence Compression

Data

Vol/set

Min TM

Rate

Max TM

Rate Rate Rate

HRT/FDT Mibits or


upper quantities for dynamical

studies

PHI science

mode 4

photospheric

context 640,00 1 10 10 10 0 0 300 300 2 60,00 204,8 204,8 204,8 209,72

both FDT &

HRT 640,00 1 0 10 0 8 8 300 300 2 52,00 177,5 177,5 181,75

lower quantities for vector

magnetometry

PHI science mode 5

PHI science mode 5 is used for global helioseismology and synoptic observations. It outputs only 2

physical quantities (continuum intensity and LOS velocity) at a varying cadence highly reduced

resolution.



Designed for Full Disk Telescope only.

Parameters

cadence, resolution (rebinning)

Resource usage

Data

processing

modes

Use case

Dataset

after

inversion


(bit/px)

Min

Cadence

Max

Cadence Compression

Data

Vol/set

Min TM

Rate

Max TM

Rate Rate Rate

FDT Mibits or


(256x256)

PHI science

mode 5 global

helioseismic/ synoptic

observations

640,00 8 10 0 10 0 0 3600 60 2 0,63 0,2 10,7 1.4-5.4 5,53

(max)

640,00 16 10 0 10 0 0 3600 60 2 0,16 0,0 2,7 1,43

(min)

(128x128) for I_c for

v_LOS

PHI science mode 6

PHI science mode 6 is PHI's default synoptic mode for daily context. It is configured to generate 1

to 4 times per day a set of all 5 or only 2 physical quantities (continuum intensity and LOS velocity)

at slightly lower resolution.

Designed for full disk telescope only

Parameters

cadence, physical quantities

Resource usage

Data

processing

modes

Use

case

Dataset

after

inversion


(bit/px)

Min

Cadence

Max

Cadence Compression

Data

Vol/set

Min TM

Rate

Max TM

Rate Rate Rate

FDT Mibits or subfield I_

c B_LOS v_LOS gamma

ph

i (s) (s) Rate (*) Mibits (kibits/s) (kibits/s) (kibits/s) (kbps)

PHI science

mode 6

daily

context 640,00 2 4 4 4 4 4 86400 21600 1 20,00 0,24 0,95 0.24-0.95

0,97

(max)

( SYNOPTIC

MODE)

only

FDT

(subfield or rebin

1024x1024)

1/day 4/day ? 0,24

(min)



OR

(see RS WG)

640,00 2 4 4 0 0 0 21600 1 8,00 0,4

PHI science mode 7

PHI science mode 7 is defined to be PHI's response to a STIX flare trigger. It is a burst or flare

mode that generates fast cadence continuum images in order to observe white-light flares.

It is designed for HRT only.

Parameters

TBD

Resource usage

Data

processing

modes

Use case

Dataset

after

inversion

Rebinning digital depth/phys. quantity (bit/px) Min

Cadence

Max

Cadence Compression

Data

Vol/set

Min TM

Rate

Max TM

Rate Rate Rate

HRT Mibits or subfield I_c B_LOS v_LOS gamma phi (s) (s) Rate (*) Mibits (kibits/s) (kibits/s) (kibits/s) (kbps)

burst/flare mode,

triggered by STIX

PHI science

mode 7 640,00 1 10 0 0 0 0 60 60 2 20,00 341,33 341,33 TBD 349,53

TBD, most likely fast

cadence continuum images



4.6 RPW

RPW is a plasma/radio wave receiver system, including high sensitivity electric and

magnetic sensors. Since the receiver system covers a very wide frequency range (near-DC to 16

MHz for electric, and 0.1 Hz to 500 kHz for magnetic), different kinds of sensors are used for the

measurements.

The RPW instrument comprises:

A Thermal Noise and High Frequency receiver (TNR-HFR) for electron measurements at the local

plasma frequency and remote detection of radio emissions. TNR-HFR will provide electric power

spectral densities from 4 kHz up to 16 MHz and magnetic power spectral densities from 10 kHz up

to 500 kHz.

A Time Domain Sampler (TDS) for waveform capture up to 500 kSPS.

A Low Frequency Receiver (LFR), covering both in-situ electric and magnetic measurements from

DC to about 10 kHz. LFR will provide both waveform and power spectral densities in this

frequency range.

These three sub-systems are connected to two different sensor units: An electric antenna unit (ANT)

and a magnetic search coil unit (SCM), both of which will be optimized to perform correctly for

near-DC as well as high frequency measurements. In particular, the antenna sensor design will be

optimized to satisfy the goal of measuring both DC/low frequency electric fields and higher

frequency radio and thermal noise emissions. A biasing unit (BIAS) will allow DC electric

measurements. The three TDS, LFR and TNR-HFR sub-systems will have a common Data

Processing Unit (DPU) that will handle commands, data and communication with S/C.

ANT: Each ANT monopole serves as a simple voltage sensor. At low frequencies, an antenna is

coupled to the local plasma potential through a photoelectron sheath. Successful measurement of

DC/low frequency plasma electric fields requires that the antenna be biased (as described below). At

sufficiently high (radio) frequencies, an antenna behaves as if in a vacuum.

Preamplifiers: Each monopole is connected to the inputs of both: A low frequency and high

frequency preamplifiers.

The LF preamplifiers will measure voltage and provide bias current, using a high input impedance

follower with a bootstrapped bias resistor and voltage source. The input stage needs to handle a high

source impedance of R ≈ 50 Mohms, and C ≈ 40 pF. Thus, it must have a low leakage current, <

10pA, for the whole mission, and a low input capacitance, < 4pF, to minimize attenuation of the

input signal. This can only be achieved by proper bootstrapping of the current generator, and by

using FET operational amplifiers.

The HF preamplifiers will provide a low noise and flat frequency response from 100Hz to more than

16 MHz. As well as the LF preamplifier, theirs input impedance will be as high as possible.



BIAS Unit: The BIAS drives a constant current to the electric antennas allowing reliable DC/LF

electric field and satellite potential measurements. The floating potential V of a conductive surface

in plasma (either an individual probe or the entire spacecraft) is determined by current balance

between the out flowing photoelectron current and the inflowing plasma electron current. In the

expected solar wind plasma, a conductive surface will charge up positively and the value of the

floating potential will depend on the density and (to a lesser extent) the temperature of the

surrounding plasma. In contrast to the spacecraft potential, the antenna potential should be anchored

tightly to the local plasma potential, and should not depend on the plasma parameters and also

should be affected as little as possible by the nearby satellite. To achieve this, we draw a constant

bias current from the antennas to the spacecraft. If the bias current is suitably chosen, then the

antennas float close to the local plasma potential and their potential is not a function of the plasma

parameters. Then measurements of the potential difference between two opposite antennas which

are anchored to the local plasma potential gives reliable measurement of DC/LF electric field

component along that direction. In addition, the potential difference between antennas and

spacecraft give an estimate of the spacecraft potential that can be used to determine the plasma

density.

SCM: The Search Coil Magnetometer SCM is an inductive magnetic sensor. It is made of a core in

a high permeability material (ferrite or perm-alloy) on which are wound a main coil with several

thousand turns and a secondary coil with a few turns. The secondary coil is used to create a flux

feedback in order to have a flat frequency response on a bandwidth centered on the resonance

frequency of the main coil. The induced voltage is raised to a proper level by a preamplifier to allow

its transportation to the analyze system.

TNR-HFR: The proposed TNR-HFR instrument is a contribution to the RPW-E experiment

consisting of a double channel radio and plasma wave spectrometer. The TNR (Thermal Noise

Receiver) is a direct conversion receiver, providing quasi-instantaneous spectra, for the electrostatic

thermal noise and/or magnetic field, while the HFR (High Frequency Receiver) is a sweeping

receiver, for the survey of high frequency radio emissions. Its analogue front end is interfaced with

three sensors, two E-field inputs (4 kHz - 16MHz) and one B-field component (10kHz to 500kHz).

TDS: The Time Domain Sampler module (TDS) will provide in-situ waveform measurements of

plasma waves around local electron plasma frequency, notably Langmuir waves found in the source

region of type II and type III solar bursts and the associated electromagnetic waves. The TDS will

perform digitization of the electric and magnetic field waveforms in the frequency range from 100

Hz to 250 kHz, their pre-processing and selection of potentially interesting events to be stored in

internal memory and later transmitted to the ground. It is known from previous observations, which

impacts of dust particles on the spacecraft body show up in the electric field data as short large

amplitude spikes of a characteristic shape. An algorithm for detection of these spikes will be

implemented in the on-board software in order to collect statistics of dust impacts as a function of

time.

LFR: The Low Frequency Receiver (LFR) performs onboard digital processing of the electric and

magnetic field data (2E, 3B, 3V) and covers a frequency range from a fraction of a Hertz to 10 kHz.

It is designed to analyze the in-situ measurements of the low frequency (f <fce) electro-magnetic



waves in the solar wind and in the extended corona. Given the limitations in the telemetry, it is

necessary to implement specific techniques to take the maximum advantage of the data. The LFR is

tailored to optimize the scientific return of the data. The LFR design gives the possibility of mixing

different types of output data, from low-level processed data (waveform data) to high-level

processed data (averaged Hermitian spectral matrices and their derived parameters), with various

data rate possibilities (continuous or cyclic transmission, adaptable frequency bandwidth as well as

adaptable frequency and time resolutions). The scientific added value stems from the choice of the

most relevant combination of the different data to be transmitted.

4.6.1 RPW observables

4.6.2 RPW modes

RPW Normal mode

Analyzer Product Data Rate

LFR Waveform 1391

BP/ASM 1469

TDS

Waveform 851

Histograms 65

Statistics 25

TNR-HFR

TNR Products 532

HFR Products 63

Spectral Power 3

RPW burst mode

Analyzer Product Data Rate (bps)

LFR Waveform 12528

BP 7056

TDS

Waveform 864

Histograms 43

Statistics 29

TNR-HFR

TNR Products 5328

HFR Products 144

Spectral Power 29



RPW SBM1

This is a triggered mode involving the dump of a rolling buffer to the SSMM on detection of an

interplanetary shock. The detailed data rates calculated here don't completely match those in the

Plasson presentation used to calculate the summary rates, although the discrepancy isn't as big as for

SBM2. This is likely based on assumptions about trigger duration used in the volume estimates in

the Plasson presentation. For both SBM1 and SBM2 the summary rates are higher so will be used

wherever possible.

Analyzer Product Data Rate (bps)

LFR Waveform 201046

BP 8861.54

TDS Waveform 33784.0

RPW SBM2

SBM2 is a triggered mode used on detection of an in situ type III radio burst. These are expected to

occur roughly once every 40 days. Note that the detailed data rates listed here (source: TM Report,

CDR issue) don't match up with the figures from the Plasson presentation used to calculate the

summary rates. This could be rounding errors or different assumptions about duration of triggers

(this is configurable).

Analyzer Product Type Data Rate (bps)

LFR Waveform 12293

BP 6849

TDS Waveform 87804



4.7 SoloHI

4.7.1 SoloHI observables

SoloHI takes visible light images of the extended corona/solar wind. The signal is made up of 4

sources: (1) photospheric light scattered by the free electrons expelled by the sun, (2) photospheric

light scattered by dust including comets and asteroids, (3) stellar, galactic and planetary sources, and

(4) instrumental stray light. Light is converted into a digital number in each pixel of the detector. An

array of pixels comprises an image. Images are compressed onboard and together with an image

header which describes the parameters of the image form a file. Each file is converted into packets

which are sent down the science channel to the SSMM. The degree of compression, including

whether there is any compression, is selected as part of the instrument schedule.

4.7.2 SoloHI modes

TM figures




SoloHI consists of 1 white-light telescope with wide FOV. The image is captured on a mosaic of

four 2048x1920 APS detectors that are read out independently. This gives flexibility for image

operations: independent exposures, cadences, etc.). Data can be read out either from the whole

detector or from selected subfields.

Typically, SoloHI images will have one of 3 typical FOVs defined below. These are used in the

observing programs/modes currently defined but could be changed in-flight if necessary.

FOV split (radial x

transverse) Downlinked pixels Typical cadence Comments

SoloHI Full frame image

(40ºx40º)

5º to 45º x 40º

***

1960 px * 1960 px

(incl. 2x2 bin)

24-36 min (inner

FOV)

30-72 min (outer FOV)

***split in 2 or 3 readout frames depending on solar distance,

e.g. 5º to 25º x 40º (inner) + 25º to 45º x 40º (outer) at

perihelion,

each with different cadence. Details in table below.

SoloHI inner FOV

subframe images

(3 images of 1.88ºx5º)

5.8º to 7.68º x 5º 13.5º to 15.38º x 5º

18.5º to 20.38º x 5º

192 px * 512 px (not binned)

96 px * 256 px

(binned 2x2)

96 px * 256 px (binned 2x2)

18-36 sec 36-72 sec

1.5-3.0 min

Subframe images typically only used at and near

perihelion (up to 0.36AU)

Radial Swath subframe

image

(40ºx5º)

5º to 25º x 5º

25º to 45º x 5º

1960 px * 256 px

(incl. 2x2 bin)

6 min (inner

FOV)

12 min (outer FOV)

Radial swath images typically only used at and near

perihelion (up to 0.36AU)




source: SoloHI UM - SSD-DOC-SOLOHI-013 Rev. B draft 5

The SoloHI baseline observing program will be defined to repeat for each occurrence of the same

unique orbit (i.e. each orbit of the trajectory within the same resonance with Venus). Therefore, a

SoloHI baseline observing program will be defined for each orbit phase in NMP/EMP and will be

executed for all orbits within that resonance.

Example of such an orbit plan is below:

Observing modes - Example plan Use case #images /

day

Science data volume /

day

SoloHI data

rate

Observing duration /

orbit

(Gbits) estimate (Kbits/s) (days) example

Perihelion programs: 0.28-0.29 AU

SoloHI Solar Wind Turbulence @perih 1296 2.22 26.5 3

SoloHI Shock Formation @perih 468 2.54 30.3 3

Near-Perihelion programs: 0.29-0.36 AU

SoloHI Near-perihelion Synoptic Program 348 1.69 20.3 5

SoloHI Solar Wind Turbulence Near-

perihelion 750 1.94 23.2 2

SoloHI Shock Formation Near-perihelion 516 2.45 29.3 2

Far-Perihelion programs: 0.36-0.42 AU

SoloHI Far-Perihelion Synoptic Program 132 1.64 19.7 7

Southern Out-of-ecliptic programs:

0.5-0.7 AU

SoloHI Southern Synoptic Program 104 0.84 10.3 8

Dependent on the trajectory

Examples of more-detailed observing program for 1 type of orbit during the mission (source:

04_130904_SoloHI_CDR_ObsProg.ppt):

http://www.rssd.esa.int/cs/livelink/3273277/SoloHI_UserManual.pdf?func=doc.Fetch&nodeid=3273277



Based on table above:

• a typical perihelion programme would produce ~25kbps (during 4 days), -> see modelled

observations HI_SHOCK_PER (DATARATE=30300 [bits/sec]), HI_TURB_PER

(DATARATE=26500 [bits/sec]),

• near-perihelion SoloHI would produce ~20kbps (during 8 days) , -> see modelled

observation HI_SYN_PER (DATARATE=20300 [bits/sec])

• ~18.5kbps even further out (during 12 days) and -> see modelled observations

HI_SYN_NEAR (DATARATE=19700 [bits/sec])

• in the far-out RS window, a data rate around 10 kbps would be reached. -> see modelled

observations HI_SYN_FAR (DATARATE=10300 [bits/sec])

(see also SoloHI concept study report Dec 2011)

Update needed:

How to organize SoloHI observations in coordination with the other instruments, i.e. does SoloHI

have 'observing modes' to choose from for each solar distance?, is still to be discussed in more

detail. Also, while the schema above may be optimal from a science perspective, the varying

downlink rate & SSMM storage limits may impose limitations on when which data rate can be used.



4.8 SPICE

4.8.1 SPICE observables

SPICE is a spectrometer that will measure plasma density and temperature, flow velocities, presence

of plasma turbulence and composition of source region plasma on the solar disk.

4.8.2 SPICE modes

TM figures


Download capacity per orbit 5.670 GB = 45.36Gbits (over 30 days)

Subtelescopes/detectors

Source: SPICE User Manual Iss 5.0 (Nov 2014)

SPICE consists of only 1 spectral telescope with slit and scanning mirror. The 2D image (1

wavelength dimension, 1 spatial dimension along the slit) is projected on 2 detectors, 1 covering

the long wavelengths, the other covering the short wavelengths.

wavelength range detector (readout) size FOV Volume/Detector readout

SW array (approx.) 70.1nm - 79.3nm

968px @0.0095nm/pix

x 800 px (along slit)

11' (14' for 30" slit) x 16' 1.55MB

LW array (approx.) 97.1nm - 105.1nm

968px @0.0083nm/pix

x 800 px (along slit)

11' (14' for 30" slit) x 16' 1.55MB

Observational modes


SPICE's primary mode of observations consist of rapid on-disk scans that characterize plasma

dynamics (SPICE Dynamics mode), alternated with slower composition scans that map the source

regions of solar wind streams (SPICE Composition Mapping). However, SPICE also has alternative

modes or 'studies' listed below. These studies will be preset on board and can be called via TC,

providing some of the study parameters (see mode pages).

Observing mode Use Case/Target Line List #

repeats

Total

Duration

Data

Rate

Max time /

orbit

(min) (Kbits/s)



SPICE Composition

Mapping

= NOMINAL MODE 1

Centre, Poles, limb,

AR

15 lines

(2 profiles, 13

intensities)

1 180 0.45

SPICE Dynamics mode

=NOMINAL MODE 2

Centre, AR, CH

H I, C III, O VI,

NeVIII profiles

+ 6 intensities

10 110 20.46

SPICE Spectral Atlas

Sun center, limb,

AR, CH Full spectrum 2 22 40.34

SPICE Limb mode

Low corona above

limb

C III, O VI, Ne VIII

profiles

+ 3 intensities

1 240 2.5

SPICE CME Watch

AR, prominence,

limb

5 spectral profiles, 10

intensities 30 22.66 hr 4.46

SPICE 30"-wide Movie

(sit&stare) center, AR 1-2 line profiles 1 10 34.44

SPICE 90"-wide Movie center, AR 1-2 line profiles 40 16 20.28

SPICE Waves mode

(sit&stare) QS, CH, AR

C III, O VI, NeVIII

profiles 5 300 50.72

SPICE Two-Exposure

mode

combi of 6 bright and

faint lines 5 300 2.86

Low Latency or Precursor programs

SPICE plans to run a mini-study before the start of each new study, that is configured exactly like

the science study but has lower resolution/cadence/slit width/... and produces ~0.1MB.

Special datasets

Observing mode Use Case/Target Line List Duration/dataset Data Rate Max time / orbit

SPICE Full Raster Scan few strong lines 32 15.1



Selection of SPICE spectral lines


Ion Wavelength (Å) Log T (K) FIP (eV) M/q

H I 1025 4.0 13.6 --

C II 1036 4.3 11.3 12.0

C III 977 4.5 11.3 6.0

O IV 787.7 5.2 13.6 5.3

O V 760 5.4 13.6 4.0

O VI 1032 5.5 13.6 3.2

S V 1037 5.5 13.6 3.2

Ne VI 786.5 5.2 10.36 8.0

Ne VII 1005 5.5 21.6 4.0

Ne VIII 973 5.6 21.6 3.3

Ne VIII 770 5.8 21.6 2.8

Mg VIII 772 5.9 7.7 3.4

Mg IX 706 6.0 7.7 3.0

Mg XI 997 6.2 7.7 2.4

Si VII 1049 5.6 8.1 4.8

Si XII 521 (2nd) 6.5 8.1 2.6

Fe X 1028 6.0 7.9 6.2

Fe XVIII 975 6.9 7.9 3.3

Fe XX 721 7.0 7.9 2.9

Auxiliary lines:


Ne VIII 780 5.8 21.6 2.8

Si XII 499 (2nd) 6.5 8.1 2.6

SPICE Composition Mapping

This is the SPICE nominal observation mode I

Target: Centre, Poles, E-W, N-S, limb, AR

Line List: 15 lines (2 profiles, 13 intensities)



Slit: 4"

Data rate calculation

Parameters are highlighted in blue.

Observing

mode

#Line

Profiles

#Px/Profi

le

#

Intensiti

es

Slit

(")

Exp.ti

me

#

Mirror

Pos

Step size Tilt

time FOV (") Data cube (px)

Durati

on

Repeat

s

Cal

c.

Dat

a

Vol

Total Data

Vol

Total

Durati

on

Data

rate

= study (max=32)

2/4/6/30

(s) =#exposures

(arcsecs) (secs) X Y X Y Z (min) (MB)

Compressed (MB)

(min) (Kbits/s)

Compositio

n Mapping 2 32 13 4" 180 60 4" 0.14

240

"

660

" 60 800 90 180 1 8.64 0.606 180.1 0.45

NeVIII,

MgIX 4' 11'

sca

n

dir

alon

g

slit

spectr

al dim

Spectral-line performance

Line-SNR and spatial resolution combinations:

line-SNR>10 at 4"x2" in C-III, O-VI, Ne-VII, and Mg-IX-AR




H I 1025 4.0 13.6 --

C II 1036 4.3 11.3 12.0

C III 977 4.5 11.3 6.0

O IV 787.7 5.2 13.6 5.3

O V 760 5.4 13.6 4.0

O VI 1032 5.5 13.6 3.2

S V 1037 5.5 13.6 3.2

Ne VI 786.5 5.2 10.36 8.0

Ne VII 1005 5.5 21.6 4.0

Ne VIII 973 5.6 21.6 3.3

Ne VIII 770 5.8 21.6 2.8

Mg VIII 772 5.9 7.7 3.4

Mg IX 706 6.0 7.7 3.0




Mg XI 997 6.2 7.7 2.4

Si VII 1049 5.6 8.1 4.8

Si XII 521 (2nd) 6.5 8.1 2.6

Fe X 1028 6.0 7.9 6.2

Fe XVIII 975 6.9 7.9 3.3

Fe XX 721 7.0 7.9 2.9

Auxiliary lines:


Ne VIII 780 5.8 21.6 2.8

Si XII 499 (2nd) 6.5 8.1 2.6

SPICE Dynamics

This is the SPICE nominal observation mode II

Target: Centre, Poles, E-W, N-S, limb, AR

Line List: 15 lines (2 profiles, 13 intensities)

Slit: 2" Step size: 2"

Data rate calculation

Observing mode

#Line Profiles

#Px/Profile

Intensities

Slit (")

Exp.time

# Mirror Pos

Step size

Tilt time

FOV (")

Data cube (px) Duration

Repeats

Calc.

Data Vol

Total Data Vol

Total

Duration

Data rate

= study (max=3

2)

2/4/6/3

0 (s)

=#exposur

es

(arcsec

s)

(secs

) X Y X Y Z (min) (MB)

Compress

ed (MB) (min)

(kbits/

s)

Dynamic

s 4 32 6 2 5 120 2 0.07

24

0

84

0 120 800 140 10.26 10

268.8

0 15.756 101.4 20.46

4 14

scan

direction

alon

g slit

spectral

dimension

for 10

repeats

for 10

repeats





line-SNR > 10 at 2"x2" in C-III, O-VI, Ne-VIII (in ARs),

and at 5"x5" in Ne-VIII (in CHs and QS)




H I 1025 4.0 13.6 --

C II 1036 4.3 11.3 12.0

C III 977 4.5 11.3 6.0

O IV 787.7 5.2 13.6 5.3

O V 760 5.4 13.6 4.0

O VI 1032 5.5 13.6 3.2

S V 1037 5.5 13.6 3.2

Ne VI 786.5 5.2 10.36 8.0

Ne VII 1005 5.5 21.6 4.0

Ne VIII 973 5.6 21.6 3.3

Ne VIII 770 5.8 21.6 2.8

Mg VIII 772 5.9 7.7 3.4

Mg IX 706 6.0 7.7 3.0

Mg XI 997 6.2 7.7 2.4

Si VII 1049 5.6 8.1 4.8

Si XII 521 (2nd) 6.5 8.1 2.6

Fe X 1028 6.0 7.9 6.2

Fe XVIII 975 6.9 7.9 3.3

Fe XX 721 7.0 7.9 2.9

Auxiliary lines:




Ne VIII 780 5.8 21.6 2.8

Si XII 499 (2nd) 6.5 8.1 2.6

SPICE Spectral Atlas

Observation configuration and data rate calculation

Observing

mode Target

Line

List

#Line

Profil

es

#Px/Pr

ofile

#Int

ens

Slit

(")

Exp.ti

me

#

Mirror

Pos

Step

size

Tilt

time

FOV

(")

Data cube

(px)

Durat

ion

Repe

ats

Cal

c.

Dat

a

Vol

Total

Compre

ssed

Data

Vol

Total

Durat

ion

Data

rate

Comm

ents

= study (max=32)

2/4/6/30

(s) =#exposures

(arcsecs)

(secs)

X Y X Y Z (min) (MB)

(MB) (min) (kbits/s)

Spectral

Atlas (A

tomic

Physics)

Sun

centre,

limb,AR

,CH

Full

spectr

um

32 64 0 4 60 10 4 5.1 40 840 10 800 2048 11.1 2 65.

54 6.666 21.5 40.34

Excepti

onal

campai

gn, but

also

used for

calibration

5s

read

out

+0.1

s

step

sc

an

dir

alo

ng

slit

sc

an

dir

alo

ng

slit

=M

AX

30'/2r

ep



line-SNR>10 at 4"x2" in C-III, O-VI, Ne-VII, and Mg-IX-AR

SPICE Limb mode


Observ

ing

mode

Tar

get

Line

List

#Line

Profil

es

#Px/Pr

ofile

#

Intensi

ties

Slit

(")

Exp.ti

me

#

Mirror

Pos

Step

size

Tilt

tim

e

FOV

(") Data cube (px)

Durat

ion

Repe

ats

Cal

c.

Dat

a

Vol

Total

Data

Vol

Total

Durat

ion

Data

rate

Comm

ents

= study (max=

32)

2/4/6/

30 (s)

=#expos

ures

(arcse

cs)

(sec

s) X Y X Y Z (min)

(M

B)

Compre

ssed

(MB)

(min) (kbit

s/s)

Limb

Low

coro

na

abov

e

limb

C III, O VI, Ne

VIII

profiles

+

3intensi

ties

3 32 3 4 60 240 4 0.1

4

96

0

84

0 240 800 102 240.6 1

39.

17 4.545 240.6 2.52



16

'

14

'

scan

direction

alo

ng slit

spectral

dimension



line-SNR >10 at 4"x2" in C-III, O-VI, Ne-VIII and Mg-IX (in ARs)

SPICE CME Watch


Obser

ving

mode

Target Line

List

#Line

Profil

es

#Px/Pr

ofile

#

Intensi

ties

Slit

(")

Exp.ti

me

#

Mirror

Pos

Step

size

Tilt

tim

e

FOV

(") Data cube (px)

Durat

ion

Repe

ats

Cal

c.

Dat

a

Vol

Total

Data

Vol

Total

Durat

ion

Data

rate

Comm

ents

= study

(max=32)

2/4/6/30

(s) =#exposures

(arcsecs)

(secs)

X Y X Y Z (min) (MB)

Compre

ssed

(MB)

(min) (kbits/s)

CME

Watch

AR,

promine

nce,

limb

5

spectra

l

profile

s, 10

intensi

ties

5 32 10 4 30 90 4 0.1

4

36

0

84

0 90 800 180 45.32 30

777.

60 45.45

1356.

2 4.46

scan

direct

ion

alo

ng

slit

spectra

l

dimens

ion



line-SNR >10 at 4"x2" in C-III, O-VI, Ne-VIII and

at 4"x4" in Mg-IX (in ARs)



Observing

mode

Tar

get

Line

List

#Line

Profil

es

#Px/Pr

ofile

#

Intensi

ties

Slit

(")

Exp.ti

me

#

Mirror

Pos

Step

size

Tilt

tim

e

FOV

(") Data cube (px)

Durat

ion

Repe

ats

Cal

c.

Dat

a Vol

Total

Data

Vol

Total

Durat

ion

Data

rate

Comm

ents

= study (max=

32)

2/4/6

/30 (s)

=#expos

ures

(arcse

cs)

(sec

s) X Y X Y Z (min)

(M

B)

Compre

ssed

(MB)

(min) (kbit

s/s)



30"-wide

Movie

(sit and

stare)

centr

e,

AR

2

line

profi

les

2 32 0 30 5 120 0 0 3

0

84

0 120 800 64 10.2 1

12.

29 2.626 10.0 34.44

To be

run very

rarely,

for

high-

cadenc

e

variabil

ity

OR

1

profi

le

extract

full slit

width

scan

direct

ion

alo

ng

slit

spectra

l

dimens

ion


N/A



Observing mode

Target

Line List

#Line Profile

s

#Px/Profile

# Intensi

ties

Slit (")

Exp.time

# Mirror

Pos

Step size

Tilt tim

e

FOV (")

Data cube (px) Duration

Repeats

Cal

c. Dat

a

Vol

Total Data

Vol

Total Durat

ion

Data rate

Comments

= study (max=

32)

2/4/6/

30 (s)

=#expos

ures

(arcse

cs)

(sec

s) X Y X Y Z (min)

(M

B)

Compre

ssed

(MB)

(min) (kbits

/s)

90"-

wide

Movie

centr

e,

AR

2

line

profi

les

2 32 0 30 5 3 28 0.9

8

8

6

84

0 3 800 64 0.435 40

12.

29 2.626 11.3 20.28

To be

run very

rarely,

for

high-

cadence

variabil

ity

OR 1 profi

le

extract full slit

width

scan direct

ion

along

slit

spectral dimensi

on


N/A

SPICE Waves mode


Observ

ing

mode

Tar

get

Line

List

#Line

Profile

s

#Px/Pr

ofile

#

Intensi

ties

Slit

(")

Exp.ti

me

#

Mirror

Pos

Step

size

Tilt

tim

e

FOV

(") Data cube (px)

Durat

ion

Repe

ats

Calc

.

Dat

a

Vol

Total

Data

Vol

Total

Durat

ion

Data

rate

Commen

ts

= study (max= 2/4/6/ (s) =#expos (arcse (sec X Y X Y Z (min) (MB Compre (min) (kbits



32) 30 ures cs) s) ) ssed

(MB)

/s)

Waves

(sit

and

stare)

QS,

CH,

AR

C III,

O

VI,

NeV

III

profi

les

3 64 0 4 5 720 0 0 4 84

0 720 800 192 60 5

109.

23 114.13 300.0 50.72

For

oscillatio

ns

measure

ments

TBC

scan direct

ion

along

slit

spectral dimens

ion


line-SNR > 10 at 4"x2" in C-III, O-VI, Ne-VIII (in ARs) and

at <4"x6" in Ne-VII (CHs and QS)

SPICE Two-Exposure mode


Observ

ing

mode

Tar

get

Lin

e

List

#Line

Profile

s

#Px/Pr

ofile

#

Intensi

ties

Slit

(")

Exp.ti

me

# Mirror

Pos

Step

size

Tilt

tim

e

FOV

(") Data cube (px)

Durat

ion

Repe

ats

Cal

c.

Dat

a

Vol

Total

Data Vol

Total

Durat

ion

Data

rate

Comm

ents

= study (max=

32)

2/4/6/

30 (s)

=#expos

ures

(arcse

cs)

(sec

s) X Y X Y Z (min)

(M

B)

Compre

ssed

(MB)

(min) (kbits

/s)

Two-

exposu

re

com

bi of

6

brig

ht

and

faint lines

3

brig

ht

line

s

3 32 0 4 5 60 4 0.1

4

24

0

84

0 60 800 96 60 5

46.

08 6.464 275.7 2.86

Profiles

are used

to

monitor

saturati

on

3

fain

t

line

s

3 32 0 4 55 60 4 0.1

4

24

0

84

0 60 800 96 5 46.08 SUM!

TBC

max

max

scan directi

on

along

slit

spectral dimensi

on

simultan

eous exposure

s!


line-SNR > 10 at 4"x2" in C-III, O-VI, Ne-VIII and Mg-IX (in ARs)



SPICE Full Raster Scan


Observi

ng

mode

Targ

et

Line

List

#Line

Profile

s

#Px/Pro

file

#

Intensit

ies

Slit

(")

Exp.ti

me

# Mirror

Pos

Step

size

Tilt

tim

e

FOV

(")

Data cube

(px)

Durati

on

Repe

ats

Cal

c.

Dat

a

Vol

Total

Data Vol

Total

Durati

on

Data

rate

Comme

nts

= study (max=

32)

2/4/6/

30 (s)

=#exposu

res

(arcse

cs)

(sec

s) X Y X Y Z (min)

(M

B)

Compres

sed (MB) (min)

(kbits

/s)

Full

raster

scan

few

stro

ng lines

3 32 0 2 4 480 2 0.07 96

0

84

0

48

0

80

0

9

6 32 1

73.7

3 32.6

Campai

gn not

listed in

EID-B

but in

RS WG doc for

first

orbits

TBC

16

ma

x

14

ma

x

(RS

WG

doc)


line-SNR > 10 at 4"x2" in C-III, O-VI, Ne-VIII (in ARs) and

at <4"x6" in Ne-VII (CHs and QS)



4.9 STIX

The Spectrometer Telescope for Imaging X rays (STIX) provides imaging spectroscopy of solar

thermal and non-thermal X-ray emissions from ~4 to 150 keV, with unprecedented sensitivity and

spatial resolution (near perihelion), and good spectral resolution.

Observationally, STIX determines the location, spectrum and timing of transient X-ray emission on

the Sun at energy ranges that encompass emission from both hot thermal plasmas and

bremsstrahlung from energetic electrons. The properties of the electrons that

generated the X-rays can be inferred from their X-ray spectrum. The distinction between a thermal

plasma and non-thermal electron population is based on the shape of the X-ray spectrum with the

latter having a characteristic power law (or broken power law) profile

and the former providing a black body spectrum (corresponding to 106 to 108 K). The spectra are

very steep and so good spectral resolution is required for their interpretation.

There is also an Iron line complex at 6.7 keV which, if isolated, can be interpreted in terms of the

thermal electron population. Since a typical flare typically generates both thermal and non-thermal

emission, which often are not co-located (for example at the top

and footpoints of magnetic loops respectively), both good spatial and spectral resolution are

required.

The observational objectives are achieved by imaging the Sun as a function of time and energy with

enough spatial, spectral and temporal resolution to match the sources of interest. Comparing the

resulting images at different energies yields the X-ray spectra of

individual features (e.g. footpoints or flaring loops). Comparing the images as a function of time

discloses the temporal behavior of the hot plasma and accelerated electrons. The data can also be

combined to yield spatially-integrated light curves and spectra. In all

cases, the basic observational datum is a single, photometrically-accurate image corresponding to a

well-defined time and energy interval.

Focusing optics are not a feasible option for arcsecond-class hard X-ray imaging within Solar

Orbiter constraints. As a result, STIX uses an indirect Fourier imaging technique based on X-ray

collimation. Conceptually, the instrument is made up of three

mechanically separate modules: X-ray transparent windows; a passive imager containing front and

rear grids; and a Detector/Electronics Module (DEM) containing passively cooled X-ray detectors

and electronics.



4.9.1 STIX observables

4.9.2 STIX modes

TM figures

Allocated TM 0.7 kbits/s

Download capacity per orbit

226.8 MB

OR 1270 MB

= 1.8144 Gbits (if only acquiring during 30 days)

= 10.16 Gbits (if extrapolated to 168 days orbit)

Subtelescopes

non

Observational modes

source:

For STIX there is only 1 data acquisition mode. Acquired science data are selected on-board

autonomously (~500bps) or selected on-ground and retrieved via TC (~100bps).

Mode Data Rate (bps) Orbital Volume (Gbits)

STIX Normal Mode - automatic onboard selection of data ~500 7.3 (if acquiring full orbit)

STIX Normal Mode - TC-requested data ~100 1.45 (if acquiring full orbit)

STIX LL ~50 0.7 (if acquiring full orbit)

STIX HK ~50 0.7 (if acquiring full orbit)

STIX has only 1 science acquisition mode, which is independent from solar activity, campaigns, etc.

Data acquisition operations onboard do change based on the incoming flux. The attenuator is

automatically used during high solar activity to suppress the count rates, cadence can be changed

too.

STIX does change the data processing and packaging onboard, based on solar activity and

automatic event detection:

During non-flaring times: STIX FPGA accumulates background counts for energy calibration.

During flares:

• STIX may autonomously change attenuator and/or enable selective pixel suppression using

predefined count rate criteria.

• STIX transmits flare flag message to s/c.

• In real time: STIX FPGA converts fine native A/D channels into detector-matched ‘science

energy channels’.

• In real time: data are accumulated as a function of science energy channel, pixel and time

bins (0.1 second or larger depending on statistics).



• In background (using pre-parameterized algorithms or in response to TC): STIX selects and

compresses data in archive memory into TM-ready packets in the 'to be transmitted’ buffer.

When requested through TC: extra TM-ready packets can be selected and transferred to the SSMM

for downlink.



4.10 SWA

SWA consists of a suite of sensors which are able to measure the three-dimensional velocity

distribution functions of the major solar wind constituents: protons, alpha particles and electrons.

The basic moments of the distributions, such as density, velocity, temperature tensor, and heat flux

vector need to be obtained under all solar wind conditions and must be sampled sufficiently rapidly

to characterize fully the fluid and kinetic state of the wind. In addition, measurements of

representative high-FIP elements (the C, N, O group) and of low-FIP elements (such as Fe, Si or

Mg) are required. These measurement challenges require an instrument suite comprising 3 distinct

sensors:

• The Electron Analyser System (EAS) to make high temporal resolution measurements of the

core, halo and ‘strahl’ electron VDFs and their moments;

• The Proton-alpha sensor (PAS) to measure the VDFs of major ion species at high time

resolution and determine their moments;

• The Heavy Ion Sensor (HIS) to measure the 3-D VDFs and determine the abundance and

charge states of prominent minor ion species.

The SWA-EAS will resolve the full 3-D velocity space distributions of solar wind electrons with

high cadence (<10 s time resolution). Since the electron thermal velocity is much higher than their

bulk velocity, even in the solar wind, this can only be achieved by a sensor, or sensors, that have a

combined field-of-view (FoV) covering a large fraction of the 4π steradians of the full sky. In

addition, the sensor is required to provide pitch angle distributions with high time resolution (ideally

at a cadence of 0.125 s). Magnetic field data from MAG is used onboard to produce the PAD’s and

is thus required at this high cadence. Note this cadence, which is slightly slower than that detailed in

the PDD, has been selected to allow synchronisation with anticipated magnetometer data rates.

The SWA-PAS sensor is capable of resolving the full 3-D VDFs of solar wind protons and alpha

particles with high cadence (<10 s time resolution), as well as measuring the bulk plasma

parameters at ultra-high time resolution of ~0.1 s to characterize the global structure and dynamics

of the 3-D inner heliosphere and improve our basic understanding of the kinetic processes and

microstate of the evolving solar wind from 0.28–1.3 AU. The PAS energy coverage and resolution,

field-of-view, angular coverage and resolution, geometric factor and time resolution are such that it

can measure solar wind protons and alpha particles for more than 99% of the time during the SO

mission profile.

The Heavy Ion Sensor (HIS) must address two fundamentally different sets of measurement

objectives. First, HIS will measure the composition and 3-D VDFs of heavy ions in the bulk solar

wind. Second, HIS will measure the composition and 3-D VDFs of the major constituents in the

suprathermal energy range. The sensor is able to resolve the full 3-D velocity distribution functions

of the prominent heavy ions at a resolution of 5 minutes in normal mode and 30 s in burst mode.

Additionally, HIS will measure 3-D VDFs of alpha particles at 4 second resolution in burst mode.

Measurements will be made up to 60 keV/e, with 64 energy steps (6-10% resolution) and 6º x 6º

resolution (15 azimuthal sectors and 6 elevation steps). The mass resolution (m/Δm) is ~5. These



will be the first such measurements in the inner heliosphere and hence any new result from HIS will

be breakthrough science.

4.10.1 SWA observables

4.10.2 SWA modes

Note that all data volumes & rates are based on powers of 10, i.e. 1Gbit = 109 bits. HK and LL are

included in normal mode.

TM Figures

Allocated TM 14.5 kbits/s

Download capacity per orbit (168 Days) 26.309 GB = 210.470Gbits

SSMM allocation

Operational Modes

Sources: SWA Budget Report Issue 1 (scaled to reach 14500 not 14848), SWA input to SAP

Planning

Mode Data Rate (bps) Duty Cycle (%) Duty Cycle (hrs / orbit) Orbital Volume (Gbits)



SWA Normal Mode 11551 99.653 4018 167.083

SWA Burst Mode (Scheduled) 413350 0.347 (min) 14 20.833

SWA Triggered Mode 447504 0.347 (max) 14 22.554

SWA's Book Keeping Algorithm trades off triggered burst modes against scheduled burst modes

and it isn't entirely clear from available documentation what the native data rates of these burst

modes are, only what proportion of the SWA TM budget is available after normal mode, LL and

HK are accounted for. Discussions with C. Owen and ISWG imply that EAS can burst for 12

minutes per day total in scheduled burst mode, and that one triggered burst mode (dump of 5 minute

rolling buffer) is equivalent to 7 minutes of scheduled burst mode. For simplicity's sake we assume

the same is true for PAS (which has many different burst mode options in any case so this could

probably reflect some form of reality). HIS does not respond to triggers.

For the purposes of calculating rates and duty cycles, then, we assume the following:

1. Scheduled burst mode involves EAS, HIS and PAS.

2. Scheduled burst mode replaces normal mode.

3. EAS and PAS respond to triggers; HIS does not.

4. Rate calculations are based on 5 minutes scheduled burst mode per day and 1 trigger

response per day. A trigger produces 7/12 of the daily burst mode volume for EAS and PAS.

Thus scheduled bursts produce 5/12 of the available daily burst mode volume for EAS and

PAS and 100% of the available burst mode volume for HIS.

SWA Normal Mode

SWA Normal mode is expected to operate for the majority of the mission, except during scheduled

bursts (~5 minutes per day). Normal Mode Provides the following data products (note that LL

products are listed here for convenience). Housekeeping represents 300bps which is included

normal mode on the summary page but not listed explicitly here

Product Cadence (s) Data Rate (bps)

EAS 3239

Electron Moments 4

3D VDFs 100

Strahl Energy Shell (LL) 100

PAS 2864

Proton moments (LL) 4

3D proton & alpha VDFs 4

8s duration proton reduced VDF snapshots 300

HIS 5148

Heavy ion rates and VDFs 300

Alpha rates and VDFs 30



Total 11551

SWA Burst Mode (Scheduled)

SWA burst mode is scheduled and coordinated with the rest of the in situ payload. Current thinking

is that there will be five minutes of 'protected' burst mode per day (on average - this will likely be

weighted towards perihelion) and a further five minutes that can be traded off against triggers (for

PAS and EAS). Scheduled burst mode replaces normal mode. PAS Burst mode is highly

configurable, trading off cadence and angular/energy coverage. Duty Cycles on the previous page

are calculated by assuming triggers occur with sufficient frequency that the second five minute

scheduled burst is always replaced.


EAS 130388

2D Electron Pitch Angle Distribution 0.125

PAS 189257

High time resolution 1D & 2D VDFs < 1

HIS 93704

Heavy ion rates and VDFs 30

Alpha rates and VDFs 3

Total 413350

SWA Triggered Mode

SWA triggered mode will definitely involve EAS. PAS involvement is subject to power constraints.

Here we assume PAS is involved. In case that it isn't there will be a proportional increase in PAS

scheduled burst mode data volume. Triggers modes do not interrupt normal mode operations, and is

not technically a mode but rather represents a dump of a 5 minute duration rolling buffer to the

SSMM.


EAS 182544

3D Electron VDFs 1

PAS 264961

3D proton & alpha VDFs 1

Total 447504





5 SCIENCE ACTIVITIES

5.1 Introduction and SOOPs

TBW

5.2 List of SOOPs

This section describes all SOOPs that have been defined up to now and are meant to cover all

science objectives. The current list is given below with all details per SOOP given in the following

sections.

• I_DEFAULT

• L_BOTH_HRES_LCAD_CH_Boundary_Expansion

• L_BOTH_LRES_MCAD_Pole-to-Pole

• L_BOTH_MRES_MCAD_Farside_Connection

• L_BOTH_MRES_MCAD_Flare_SEPs

• L_FULL_HRES_HCAD_Coronal_Dynamics

• L_FULL_HRES_HCAD_Eruption_Watch

• L_FULL_HRES_LCAD_MagnFieldConfig

• L_FULL_HRES_MCAD_Coronal_He_Abundance

• L_FULL_LRES_MCAD_Coronal_Synoptic

• L_FULL_LRES_MCAD_ProbeQuadrature

• L_FULL_MRES_MCAD_CME_SEPs

• L_IS_SoloHI_STIX

• L_IS_STIX

• L_SMALL_HRES_HCAD_Fast_Wind

• L_SMALL_HRES_HCAD_SlowWindConnection

• L_SMALL_MRES_MCAD_Ballistic-connection

• L_SMALL_MRES_MCAD_Connection_Mosaic

• R_FULL_HRES_HCAD_Density_Fluctuations

• R_FULL_LRES_HCAD_GlobalHelioseismology

• R_SMALL_HRES_HCAD_AR_Dynamics

• R_SMALL_HRES_HCAD_Ephemeral

• R_SMALL_HRES_HCAD_PDF_Mosaic

• R_SMALL_HRES_HCAD_Photospheric_Dynamics_Structure

• R_SMALL_HRES_HCAD_RSburst

• R_SMALL_HRES_HCAD_WaveStereoscopy

• R_SMALL_HRES_LCAD_Composition_vs_Height

• R_SMALL_HRES_LCAD_FineScaleStructure

• R_SMALL_HRES_MCAD_PolarObservations

• R_SMALL_MRES_MCAD_AR_LongTerm



5.2.1 I_DEFAULT

Standard In situ operations out of remote sensing windows. Will contribute to all in situ only

objectives. Coordinated bursts schedules to obtain equal coverage in heliocentric distance rather

than time.

Instrument Mode Comment

EPD Normal + Burst

MAG Normal + Burst

RPW Normal + Burst Triggers Active

SWA Normal + Burst

SAP

objective Target Duration

Opportunity(e.g.,

orbital

requirements,

solar cycle phase,

quadrature ...)

Operationa

l

constraints

Additional comments

2.3.2.2

Identify

interplanetary

shocks and

characterise

their spatial

and temporal

evolution

In Situ

Sufficient

Coverage

for good

statistics

Multi

spacecraftstudy(SPP

) in multiple

orientations (radial,

spiral, quadrature

alignments) –

should come for

free.

For quadrature, one

s/c would need RS,

and the other in-situ,

so again for free if

WISPR R< 0.5 AU

EMC Quiet

Burst modes

scheduled for flat

radial coverage



2.3.3 Resolve

the

interplanetary

shock field

and plasma

structure

down to the

spatial and

temporal

scales

comparable

and smaller

than the

typical ion

scales.

In Situ

Sufficient

Coverage

for good

statistics

EMC Quiet

Burst

Modes,mostlytriggere

d.

SPP Good to have.

2.3.4 Shock-

surfing

acceleration

mechanism

In Situ

Heliosphere

Sufficient

Coverage

for good

statistics

Radial dependence EMC Quiet

Normal + Burst modes

Triggered by RPW

(this is important to

know for low

telemetry periods

when the triggering

would be de-

activated).

SPP Good to have.

2.3.5

Understand

the radio

emissions

from the

ICME driven

shocks

In SItu

Sufficient

Coverage

for good

statistics

Good coverage of

different radii EMC Quiet

RPW Triggers

important.

1.2.2.7 Study

the

correlation

degree

between

velocity and

magnetic

field

fluctuations

in the

interplanetary

space

In Situ

Sufficient

Coverage

for good

statistics

Good coverage of

different radii and

latitudes

EMC Quiet Normal Mode

Sufficient



1.1.2.6

Disentangle

the spatial

and temporal

variability of

the slow wind

In Situ

1.1.2.8

Determine

the velocity,

acceleration

profile and

the mass of

the transient

slow wind

flows

In Situ

1.1.4.1.3

Identify

reconnection

exhausts in

the solar wind

In Situ

1.1.4.1.5

Identify and

characterise

the solar wind

reconnection

physics in

current sheets

with

thickness

down to the

ion scales and

smaller

1.2.2.1

Determine

where energy

is deposited

in the solar

wind



1.2.2.2 What

drives the

evolution of

the solar wind

distribution

functions in

situ?

EMC Quiet

• Long-term

observations in

normal mode.

• Good radial

spread of short

duration bursts

over the

mission.

• Radial

coverage

including

perihelion.

• SO-SPP(-

Earth) radial

alignments.

1.2.2.4

Identify and

characterise

the solar wind

reconnection

physics in

current sheets

with

thickness

down to the

ion scales and

smaller.

3.1.1.2.2

Composition

variations

In Situ

Seed

population

specification

from the

heavy ion

composition

of solar wind

and

suprathermal

s in the inner

heliosphere

timing

Statistics

Normal mode

sufficient.

SWA/HIS

EPD/SIS

RPW/MAG not

needed.

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/3.1.1.2.2+Composition+variations





3.1.1.3 How

are superhalo

particles

accelerated

continuously

in the corona

and solar

wind?

In Situ

Particles

acceleration

Statistics Good radial spread.

However, SPICE

could be used to

observe non-thermal

electron distributions

at the limb, e.g in

combination with SPP

in quadrature.

However, if SPICE

cannot observe the

superhalo energies,

then it is not needed

here (TBC).

3.1.1.5 Do

proton-

amplified

Alfvén waves

play a role in

accelerating

particles at

shocks?

In Situ

Alfvén

waves role in

SEP

acceleration

Statistics

Perihelion SWA Burst mode

3.1.2.0.1

Measure the

enhancements

of trans-iron

elements in

impulsive

SEPs (to be

deleted -

trans-iron

SEPs cannot

be detected

by EPD)

In Situ Statistics

EPD (SIS) Normal

mode sufficient

SWA/MAG/RPW not

needed

RS window not

necessarily needed (for

this specific

objective), however it

would be great if we

can identify the solar

source as for all

Objective 3 (SPICE

would be nice to have)

It seems to be feasible

with SIS.











3.1.3.1 How

shocks can

accelerate

electrons to

relativistic

energies

(never

observed for

shocks near 1

AU)?

In Situ Statistics

Perihelion Normal + Burst modes

3.2.4 How do

large and

small-scale

structures

modulate

particle

fluxes?

In Situ

Corona +

Heliosphere

Long term

observation

s

Radial dependence

RS context is good

to have (not

necessary).

Good to

have Metis

compatible

May need full-disk RS

+ Metis for context.

SWA regular

scheduled and

triggered burst by

RPW (Multiple bursts

per day only needed

for the small-scale

structures but we

cannot predict when

this will happen in

order to plan for

multiple scheduled

bursts per day, except

if we know we are in a

particularly active

period. Large scale

CIRs/ICMEs may

rather need triggered

burst)

3.2.0 What

controls the

escape of the

particles to

the

heliosphere?

In Situ

Heliosphere Statistics Perihelion

EPD

MAG

RPW

1.1.3.2.3

How does the

heliospheric

magnetic

field

disconnect

from the Sun?

In Situ

Heliosphere

Radial dependence

& Perihelion










3.2.1.2

Measurement

s of SEP

events time

profiles and

anisotropy in

order to probe

solar wind

turbulence

In Situ

Heliosphere Statistics

Close to perihelion

and different range

of distances

Needs to be done for a

range of distances to

monitor the SEP

contributions from

solar source and IP

turbulence.

Also partly addressed

by L_IS_STIX and

better by

L_IS_SoloHI_STIX.

3.2.1.3

Identify

dropouts and

measure

scattering of

SEPs by

turbulence.

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/L_IS_STIX



5.2.2 L_IS_STIX

This SOOP is aimed at understanding X-Ray emission from energetic particles. STIX is leading this

SOOP, while IS payload provides continuous observations.

Default SOOP duration: currently modelled to run during all RS windows

Pointing requirements: disk centre preferred but not mandatory

Triggers: enabled


STIX (leads) STIX Normal Mode

SoloHI N/A

Metis N/A

EUI N/A

PHI N/A

SPICE N/A

MAG MAG Normal Mode + MAG Burst Mode

EPD EPD Normal Mode + EPD Burst Mode

RPW RPW Normal Mode + Burst modes Triggers on

SWA SWA Normal Mode and SWA Burst Mode (Scheduled)

SAP objective Target Duration

Opportunity

(e.g., orbital

requirements,

solar cycle

phase,

quadrature ...)

Operational

constraints Additional comments

3.1.1.4 How

can SEPs be

accelerated to

high energies so

rapidly?

Corona +

Heliosphere

Particles

acceleration

Statistics

Perihelion

SWA Burst mode








3.1.2.3 Flare

seed particles.

Corona &

Heliosphere

SEP events

with large

initial Fe/O

ratios

Particles

accelerated

by flares

Statistics

Close to

perihelion

3.1.2.7 How are

so many

electrons

accelerated on

such short time

scales to

explain the

observed hard

X-ray fluxes?

Corona &

Heliosphere

X-ray

signatures of

energetic

particle

interactions

at loop

footpoints,

or on loops

themselves

Statistics

Perihelion (if

possible, but to

be connected to

earth is more

important)

Best when the

solar limb from

SO is connected

to Earth, or

other s/c (not

necessarily the

limb, but behind

the limb, up to

20 degrees)

To use RS context from

Earth

STIX: High-cadence

energy resolved imaging

3.2.1.2

Measurements

of SEP events

time profiles

and anisotropy

in order to

probe solar

wind turbulence


Close to

perihelion and

different range

of distances

Needs to be done for a

range of distances to

monitor the SEP

contributions from solar

source and IP turbulence


by I_DEFAULT or better

with L_IS_SoloHI_STIX.

3.2.1.3 Identify

dropouts and

measure

scattering of

SEPs by

turbulence.






5.2.3 L_IS_SoloHI_STIX

This SOOP is aimed at measuring variability in SEP events trough corona and heliosphere. IS

instruments provide continuous observations while SoloHI gives spatial context.

Default SOOP duration: 10 days

Pointing requirements: disk centre

Triggers: enabled


SoloHI

SoloHI Shock or

Turbulence and

Synoptic modes

The distance to the sun will define the SoloHI mode: shock

and turbulence mode, in combi with HI_SYN_PER to be

scheduled only within 0.4AU, further out only synoptic modes

(HI_SYN_NEAR or HI_SYN_FAR).

EUI N/A

PHI N/A

STIX STIX Normal Mode Needed only for 3.2.1.2

SPICE N/A

Metis N/A

EPD Triggered Burst Mode

MAG MAG Normal Mode

and MAG Burst Mode

RPW RPW Normal Mode

and RPW Burst Mode

SWA

SWA Triggered Mode

or SWA Burst Mode

(Scheduled)

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/SoloHI

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/STIX+Normal+Mode



SAP objective Target Duratio

n

Opportunity

(e.g., orbital

requirement

s, solar cycle

phase,

quadrature

...)

Operationa

l

constraints

Additional comments

3.1.1.2.3

Warped shock

fronts

Corona +

Heliosphere

Gradual SEP

events

Statistics

Good radial

spread.

Multispacecraft

Multiview point data preferred

to address this objective

IS Triggered burst

3.1.1.2.4

Turbulence and

inhomogeneitie

s

Magnetic

field, plasma

wave and

solar wind

measurement

s to

determine

turbulence

levels and

identify

shock

passages

Statistics

Good radial

spread.

Near

Perihelion

(EPD)

EMC

Quiet

SoloHI is used with

Solarprobe (turbulence mode).

Burst Modes: scheduled or

triggered.

MAG: high-cadence magnetic

field

RPW: high-cadence electric

and magnetic field, power

spectral densities

3.2.1.2

Measurements

of SEP events

time profiles

and anisotropy

in order to

probe solar

wind

turbulence

Heliosphere

Statistics

Close to

perihelion

and different

range of

distances

Needs to be done for a range

of distances to monitor the

SEP contributions from solar

source and IP turbulence

SoloHI Turbulence mode

STIX


by I_DEFAULT or L_IS_STI

X.

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/3.2.1.2+Measurements+of+SEP+events+time+profiles+and+anisotropy+in+order+to+probe+solar+wind+turbulence














3.2.3 Properties

of near-Sun

shocks, their

fluctuations

and particle

acceleration

Corona +

Heliosphere Statistics Perihelion EMC Quiet

MAG burst mode (good

trigger needed)

SWA burst mode (triggered)

RPW

SoloHI Shock mode

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/3.2.3+Properties+of+near-Sun+shocks%2C+their+fluctuations+and+particle+acceleration








5.2.4 L_FULL_LRES_MCAD_Coronal_Synoptic

Synoptic SOOP designed for CME science and global coronal structure. Typically runs for a whole

RS window.



Triggers: IS + Metis triggers active


EUI FSI Synoptic mode (S) 10 min cadence

Metis

METIS standard modes: GLOBAL mode,

METIS special modes CMEOBS whenever triggered

(model as max 3 CMEs per RSW)

CME Watch On

PHI FDT Synoptics: PHI science mode 6 with 6hr cadence 6 hour cadence

SoloHI Normal Operations: SYN + SHOCK (at perihelion)

(modelled as HI_SYN_NEAR for now)

SPICE N/A

STIX STIX Normal Mode

EPD Normal + Burst Mode

MAG Normal + Burst Mode

RPW Detection Mode Burst Triggers Active

SWA Normal + Burst Mode




Opportunity

(e.g., orbital

requirements,

solar cycle phase,

quadrature ...)

Operational


2.1.1.2 CME

structure

Full Disk Long

Duration

Quadrature ideal,

perihelion would

be best, but useful

at all distances.

No Offpoints

For EUI could do

onboard prioritisation

to increase cadence -

discard 9/10 images

unless there's a trigger.

2.1.1.3 CME

evolution

Full Disk Long

Duration

Higher latitude

orbits particularly

interesting.

Angular Separation

from other

spacecraft is a

bonus

No Offpoints Temperature evolution

from Metis in UV

1.1.3.2 How

does the Sun's

magnetic field

link into space?

Corona +

Heliosphere

Metis

compatible

Coordination with SPP

is a bonus

1.3.3 Plasma

turbulence

variability

Full Disk Long

Duration

Perihelion good for

SoloHI

contribution

Several Latitudes

No Offpoints

EMC Quiet

Radial alignment with

SPP useful

1.3.4 Plasma

turbulence

anisotropy

Full Disk Long

Duration

Radial

Dependence,

EMC Quiet

No Offpoints

SPP important to have

1.2.1.10

Heating in

flaring loops vs

heating in

active regions

Full Disk Long

Duration

More statistical

study: having

STIX on,

observing all the

flares and using

EUI synoptics to

find out about the

source region

N/A







3.1.2.2

Evaluate how

significantly

large flares

contribute

directly to

gradual SEP

events

Corona &


Near perihelion in

order to get 'as-

pure-as -possible'

signal

Metis

compatible

FSI 10 min cadence

(default)

SoloHI, Metis, STIX

Low resolution does

not seem to be a

problem

3.1.4.1 Hard X-

ray emission of

escaping

electron beams

(thin-target

emission)

Corona

Statistics

Near perihelion to

get best spatial

resolution with

STIX

RPW

STIX

3.1.4.2 X-ray

emission from

electrons

accelerated at

CME shocks

Corona &

Heliosphere

Statistics

NOT at Perihelion Metis

compatible

RPW

STIX

Metis

We cannot hunt for

this event (too specific

requirements: CME

lifting off from behind

the limb to shield the

X-rays from AR

footpoints, to allow

STIX to observe hard

X-rays from shock)

and we hope we get it

for free through this

SOOP

4.2 What are

the properties

of the magnetic

field at high

solar latitudes?

[source: SOL-

PHI-MPS-

MN5100-TN-2]

Full disk

synoptic

during

solar

minimum

opposition with

NEO

Co-

observations

with NEO

global field

maps

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/3.1.2.2+Evaluate+how+significantly+large+flares+contribute+directly+to+gradual+SEP+events










5.2.5 L_FULL_LRES_MCAD_ProbeQuadrature


Pointing requirements: Disk Centre

Triggers: only IS triggers enabled

SOOP designed to study the corona while Solar Probe Plus is in quadrature with Solar Orbtier, PHI

& EUI provide context, Metis and SoloHI provide imagery of solar wind that will be encountered

by solar Probe Plus.


PHI PHI science mode 6 (FDT) Full FoV, 6 hours cadence, 5 quantities

EUI FSI Synoptic mode (S)

Metis METIS special modes (PROBE mode) 1-10 minute cadence depending on

heliocentric distace

SoloHI SoloHI mode depending on heliocentric

distance


EPD Normal Mode

MAG Normal Mode


SWA Normal Mode


Opportunity

(e.g., orbital

requirements, solar

cycle phase,

quadrature ...)

Operational

constraints

Additional

comments

1.2.2 What

mechanisms heat

and accelerate the

solar wind?

helmet

streamers

several hours

METIS at

high cadence

SPP in quadrature

with Solar Orbiter

METIS co-

observations

required





5.2.6 L_FULL_MRES_MCAD_CME_SEPs

This SOOP is aimed at measuring in-situ SEPs properties and kinetics and linking them to the

observing SEPs CME acceleration. Location, timing, and motion of CMEs and shocks.

SoloHI and Metis are leading this SOOP, while IS payload provides continuous observations. Disk

centre pointing needed.



Triggers: IS+Metis triggers active


SoloHI

(leads) SoloHI Shock and Synoptic modes

Metis (leads) METIS standard modes GLOBAL, METIS special modes

CMEOBS

Disk centre pointing

needed

EUI FSI Synoptic mode (S)

FSI 10 min cadence

(default)

PHI N/A

SPICE N/A


EPD EPD Normal Mode

MAG MAG Normal Mode

RPW RPW Normal Mode

SWA SWA Triggered Mode or SWA Burst Mode (Scheduled)





Opportunity

(e.g., orbital

requirements, solar

cycle phase,

quadrature ...)

Operational

constraints

Additional

comments

3.1.1.1 Where

and when are

shocks most

efficient in

accelerating

particles?

Corona +

Heliosphere

Particles

acceleration

Statistics

Good radial spread.

Good to have

Metis

compatible

SWA 5 minutes

captures

SoloHI needed for

estimating the total

CME energy

FSI not needed

3.1.1.2.1

Intensity

variability

Heliosphere

Gradual

SEP events

Statistics

Preferable with Earth

images and SO +/- 70

degrees from Earth-

Sun line (probably

because gradual

events extend up to

100 degrees in

longitude)

Metis

compatible

MAG/RPW/SWA

not needed

FSI 5 min cadence

for source

identification (FSI is

10 min cadence by

default)

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/3.1.1.2.1+Intensity+variability





3.1.1.7 Are there favourable environments for

particle acceleration?

Corona +


FSI 10 min cadence

(default)





5.2.7 L_FULL_HRES_LCAD_MagnFieldConfig

High spatial resolution synoptic operations with full disk RS instruments, Metis and insitu

instruments. Low cadence (1/hour to 1/day). Most relevant for windows or parts of windows with

very restricted TM and for out of window synoptic campaigns if these are possible.

Default SOOP duration: 1 day


Triggers: only IS triggers active


PHI PHI science mode 6: FDT at highest spatial resolution

(2Kx2K), cadence at least 1 per day

model cadence as twice per

day

EUI FSI Reference Synoptic mode (R), highest spatial

resolution (3Kx3K), low cadence (hrs to 1/day):

Configure instrument to

enhance off-limb structures

model cadence as 1/day

Metis

synoptic programme for magnetic field structure

(METIS standard modes: GLOBAL and/or LT-

CONFIG)

no reaction to CMEs needed

EPD normal

MAG normal

SWA normal

RPW normal




Opportunity

(e.g., orbital

requirements, solar

cycle phase,

quadrature ...)

Operational

constraints

Additional

comments

1.1.3.2.1 How

does the Sun's

magnetic field

change over

time?

Full disk,

photosphere

and corona

long-term

observations

(RS synoptics

in between

RSW?)

• far side

• from Earth

alignment to

far side

(different

separation

angles with

Earth for

projection

effects)

• range of

latitudes,

definitely

repeat at

highest

latitudes late in

mission

RSW

placement if

no synoptics

allowed

EMC quiet for

linkage

science

1.1.3.2.2 How is

the heliospheric

current sheet

(HCS) related to

coronal

structure?



5.2.8 L_FULL_HRES_MCAD_Coronal_He_Abundance

Simultaneous observations of the resonantly scattered component of He+ emission by EUI/FSI 30.4

nm, and neutral hydrogen by Metis Lyα (121.6 nm) to examine helium abundances in the corona /

inner solar wind. These can be compared with SWA.

Compared to its value in the solar convective envelope, the helium abundance in the in-situ

measurements of the fast and slow solar wind has long been known to be depleted relative to

hydrogen, with occasional transient exceptions (Bochsler 1998, SSRv 85, 291). In the slow solar

wind, the degree of depletion has more recently been shown to depend upon the wind speed and the

level of solar activity (Aellig et al. 2001, GeoRL 28, 2767). Measurements of the helium abundance

in the corona, associated to measurements of the coronal outflow velocity, will provide evidence for

the degree of correlation between wind speed and helium abundance and allow identification of the

source regions of the slow wind streams with different helium abundance. During quasi-corotation

the intrinsic evolution of magnetic topology will be observed and thus its influence on the wind

parameters (such wind outflow velocity and helium abundance) will be directly assessed. The

abundance can be derived from simultaneous observations of the resonantly scattered component of

singly ionized helium by EUI/FSI in its 30.4 nm channel and of that of neutral hydrogen by Metis in

Lyα (121.6 nm).

Useful contributions can be given by SPICE, mapping the near-surface elemental abundances,

including that of helium (TBC), which constitutes a reference for establishing abundance variations

in the wind. PHI can also contribute, providing data suitable for coronal magnetic field

extrapolations.

Note: this is essentially an EUI/Metis sub-objective, but in-situ may be interested also. For instance,

SWA/HIS will measure the α/p density ratio.


EUI FSI Synoptic mode (S), 20-min

cadence

Occultor likely to be used in conjunction with this

mode

Metis

MAGTOP, 20-min cadence,

duration ≥ 2 hours

WIND, 20-min cadence,

duration ≥ 2 hours

Global maps of:

• neutral hydrogen Lyα intensity

• electron density

• outflow velocity in corona

PHI PHI science mode 6: FDT, 6-

hour cadence Data suitable for coronal magnetic field extrapolation



SoloHI Measure solar wind speed above potential source

region in co-ordination with Metis

SPICE SPICE Composition Mapping

STIX

EPD Normal Mode

MAG Normal Mode


SWA Normal Mode


Opportunity

(e.g., orbital

requirements, solar

cycle phase, quadrature

...)

Operational

constraints

Additional

comments

1.1.2 Source

regions of the

slow solar

wind

1.1.2.12

Abundance of

helium as a

function of

height and

latitude in the

corona as a

tracer of the

source regions

of the slow

solar wind

Inner

corona

within

EUI/FSI

and

Metis

FOVs

few hours

per day

• Inside 0.45–0.5

AU (optimised

distance for

EUI/FSI's

occulter)

• or perihelion, for

quasi-corotation

measurements.

Disc-centre

pointing

Earth view

beneficial before

RSW to estimate the

global solar

magnetic field. Is

this because another

view than PHI/FDT

is needed, e.g., to get

front- and back-side

measurements?





5.2.9 L_FULL_HRES_HCAD_Eruption_Watch

Full-Disk, high resolution SOOP designed to catch eruptive events.



Triggers: enabled for both IS and RS


EUI

Global Eruptive Event Mode:

FSI Global eruptive event

mode (G)

Triggered:

EUI will be in global mode all the time (i.e. full SOOP

length) but it will only prioritize the data of 1-2 events.

Global mode generates 4,42 Gbit/hr. Let's say for now

that we flush 2 hrs of data = ~8Gbits.

Metis

METIS standard modes:

GLOBAL + CMEOBS on

trigger

(modelled as 2 CME events of

1 hr in the 1-day-SOOP)

CME Trigger on

PHI

FDT, 2-5 minute cadence,

highest spatial resolution

(PHI science mode 4 with

FDT, default is 5 mins

cadence)

Selection of data so will check LLD

SoloHI Combination of shock and

synoptics at perihelion

Combine HI_SHOCK_PER + HI_SYN_PER (each 50%

of time)

SPICE SPICE Waves mode

Could be SPICE Waves mode and SPICE Composition

Mapping interleaved if off-pointing possible.

In this case, we would use observation called

SPICE_WAVES_COMP in modelling. But

SPICE_WAVES will do for SAP v.0

STIX Standard Operations: STIX

Normal Mode

MAG Normal Mode

EPD Normal Mode

RPW Detection Mode Detection algorithms active



SWA Normal Mode



SAP


Opportunity

(e.g., orbital

requirements,

solar cycle phase,

quadrature ...)

Operational


2.1.1.1

CME

initiation

Full

Disk

1 Day

(limited by

EUI/PHI

internal

memory)

Perihelion

preferred.

Quadrature with

Earth preferred.

Interesting

throughout the solar

cycle.

No

offpointing

beyond Metis

limit.

EUI/PHI at highest cadence

appropriate to spatial

resolution (could be slower if

further away). SPICE Sit &

Stare in waves mode to try and

catch EUV waves.

Metis Modes:

• GLOBAL (before the

event, if possible), min.

obs time 2 hr, data

volume ≤ 300 Mb.

• CMEOBS, starts after

CME flag rise, min. obs

time 1 hr (high

cadence, 1 min), data

volume ~ 2.137 Gb.

• GLOBAL (after the

event), min. obs time 2

hr, data volume ≤ 300

Mb.



5.2.10 L_FULL_HRES_HCAD_Coronal_Dynamics

This SOOP is aimed at observing structures in the outer corona and linking them to the heliosphere

observed in-situ. Metis and SoloHI are leading this SOOP, while IS payload provides continuous

observations. Synoptic support from other full disk RS instruments. Disk centre pointing preferred.




See SOOP 5 defined for SOWG8 planning exercise.

(SOOP will be modelled with 1 day duration, can be repeated as many times as needed)


SoloHI

(leads)

SoloHI combination of high cadence TURB and

synoptic mode (model as HI_SYN_NEAR for now)

Metis

(leads)

Generic program like WIND (METIS standard

modes) interleaved with FLUCTS (METIS special

modes) (FLUCTS runs 1 hr/day)

disk centre pointing preferred

PHI FDT synoptic: PHI science mode 6

PHI may be processing in

between observations. all data

gets downloaded

EUI FSI synoptic: FSI Synoptic mode (S) all generated data gets

downloaded

EPD Normal Mode

MAG Normal Mode


SWA Normal Mode



SPICE

(optional)

SPICE currently proposes to use either SPICE Limb

mode or SPICE Dynamics, depending on

requirement for Metis-compatibility

Limb target (not Metis-Compatible; SPICE Limb

mode) :

• Slit: 4"

• Exposure time: 5 s

• X positions: 224

• FoV: 15' x 11'

• Nº repeats: 10

• Observation time:

o 18 mins per study

o 3.2 hours total.

Active region, if before or after Metis observatiosn

(SPICE Dynamics) :

• Slit: 4"



• FoV: 8.5' x 11'

• Nº repeats: 1?

• Observation time: 2.1 hours

• Limb active region target

best if present. Cannot

participate if Metis is

observing

Lines for SPICE Limb mode:

• H I 1025 Å,

• C III 977 Å,

• O VI 1032 Å,

• Ne VIII 770 Å,

• Mg IX 706 Å,

• Si XII 520 Å (2nd order)

– 3 profiles and 3 intensities.

Lines for SPICE Dynamics:

• 4 profiles and 6

intensities.




Opportunity

(e.g., orbital

requirements, solar

cycle phase,

quadrature ...)

Operational

constraints

Additional

comments

1.1.2.7 Trace streamer

blobs and other

structures through the

outer corona and the

heliosphere.

Corona +

Heliosphere

A few

Days

Quadrature with

Earth for combining

SO RS with L1 IS

and Earth

Coronagraphs with

SO IS

Metis

compatible

Quadrature or

radial

alignment with

SPP would be a

bonus

1.3.2 How is turbulent

energy dissipated and

how does turbulence

evolve within the

heliosphere?

Corona +


Radial Alignment

with anything else.

Metis

compatible

Radial

alignment with

SPP would be a

bonus

1.1.2.5 Structure and

evolution of streamers

Corona +


Quadrature with

Earth for combining

SO RS with L1 IS

and Earth

Coronagraphs with

SO IS

Metis

compatible

SPICE

participates if

possible.

2.3.1 Coronal shocks

Corona +


Near perihelion for

highest spatial

resolution and best

spatial coverage in

the corona

Earth side for radio

obs from ground and

magnetic field

models for help with

post facto analysis

Metis

compatible

Quadrature

with SPP would

be a bonus



2.3.2.1 Understand

coronal conditions

under which the

shocks form and

determine the

interplanetary

conditions where they

evolve

Corona +


Near perihelion for

highest spatial

resolution and best

spatial coverage in

the corona

Earth side for radio

obs from ground and

magnetic field

models for help with

post facto analysis

Metis

compatible

Quadrature

with SPP would

be a bonus

2.3.2.3 Study heating

and dissipation

mechanisms at shocks

with radial distance

Corona +

Heliosphere Statistics Close to perihelion

Metis

compatible

EMC quiet

Burst modes

most important

here

Metis needs to

see within 5RS

for Lyα

2.3.2.4 Identify

mechanisms that heat

the thermal solar wind

particle populations

near shocks and

determine their

energy partition

Corona +


Good alignment

with SPP (Radial or

Quadrature) would

be beneficial

EMC quiet

Metis not

strictly needed

for this one but

could still

provide useful

context.



5.2.11 L_SMALL_MRES_MCAD_Ballistic-connection

In this SOOP, the spacecraft points at the modelled ballistic connection point, so this involves

tracking.


Pointing requirements: offpointing


Based on SOOP 4 b as planned during PlanningExercise2016_withResults.pptx


EUI

FSI Synoptic mode (S) with FSI

HRI in EUV & LYA Coronal hole

modes (C)

FSI in synoptic

HRI in CH mode at cadence ~ 900s;

PHI

PHI/HRT synoptic program:

HRT in PHI science mode 2 at

900-s cadence

ideally regular flushes;

SoloHI SoloHI: Nominal synoptic

perihelion program Model as HI_SYN_NEAR

SPICE SPICE composition & dynamics

interleaved

Use observation called SPICE_WIND_CONNECT

in modelling.


EPD Normal Mode + regular burst

mode Close mode until at least the end of RSW

MAG Normal Mode + regular burst

mode

RPW Normal Mode + regular burst

mode Burst Triggers Active, selective OK.

SWA Normal Mode + regular burst

mode

SAP


Opportunity

(e.g., orbital

requirements,

solar cycle

phase,

Operational




quadrature ...)

1.1.2.7 Trace

streamer

blobs and

other

structures

through the

outer corona

and the

heliosphere.

Near-

quadrature, so

that SoloHI

can image

Earth-directed

blobs

objective also addressed by

L_FULL_HRES_HCAD_Coronal_Dynamics



5.2.12 L_SMALL_MRES_MCAD_Connection_Mosaic

This SOOP is designed to let the high-resolution RS observations cover a wider area than normal, in

particular the SPICE FOV. This SOOP is to be used in particular with a mosaic of pointings, though

it does not necessarily need to be. Alternatively it could be used when we point for a limited amount

of time (few hours typically) to the most likely connectivity point, within a mainly sun-disk-centred

observation.

Default SOOP duration: 3 hours

Pointing requirements: mosaic of offpointing


See SOOP 4a defined for SOWG8 planning exercise.




SPICE

(leads)

SPICE Composition Mapping (30mins

per map)

• Slit: 4’’or 6” • Exposure time: 180 s

• number of X positions = 10

(inferred, to maintain 30 minute

cadence)

• Field of View: 1.2’ x 11’

o depending on step size

chosen – i.e., not the

same as the slit width –

this could cover a

larger region.

• Number of repetitions of the

study: 6 (one at each position)


Mosaic of 6 positions would take bit more than 3 hours

Lines for SPICE Composition Mapping (from 1.1.2.2

Does slow and intermediate solar wind originate from

coronal loops outside of coronal holes?) :

– 15 lines (2 profiles+ 13 intensities)

• Ne VIII 770 Å,

• Ne VIII 780 Å,

• Mg IX 706 Å,

• O II 718 Å,

• O IV 787 Å,

• O V 760 Å,

• O V 761 Å,

• O VI 1032 Å,

• Ne VI 999 Å,

• Ne VI 1010 Å,

• Mg VIII 772 Å,

• Mg VIII 782 Å,

• C III 977 Å,

• Fe III 1017 Å

• Si II 992 Å

N.B. This line list for SPICE Composition Mapping is

similar to that of

L_BOTH_HRES_LCAD_CH_Boundary_Expansion,

except for O VI 1037 (removed) and Si II 992 (added).

PHI PHI science mode 2 with HRT, 15 mins

cadence 2 PHI/HRT datasets per SPICE map

EUI

EUV & LYA Coronal hole modes (C)

with HRI, 10-15 mins cadence


2-3 HRI images per SPICE map

SoloHI

SoloHI Synoptic program

modelled as HI_SYN_NEAR

cadences adapted to applicable sun distance and TM

corridor




Metis

will generally be SAFE+door closed,

due to off-pointing mosaic.

IF far enough from Sun: METIS standard modes for context for

connection

EPD Normal Mode

MAG Normal Mode


SWA Normal Mode



SAP


Opportunity

(e.g., orbital

requirements, solar

cycle phase,

quadrature ...)

Operational

constraints

Additional

comments

1.1.2.7 Trace

streamer blobs

and other

structures

through the

outer corona

and the

heliosphere.

N-S mosaic

centred on

most likely

connection

point

3 hrs +

slews for

mosaic of

6 positions

(30mins

dwell

time)

• RS window -

quadrature

with Earth:

SoloHI

images Earth-

directed

blobs, Earth

images blobs

that will hit

Solar Orbiter

• Daily N-S

mosaic driven

by SPICE

composition

maps: 6

positions, 30

mins dwell time

at each position

• Pre-window

observations

with EUI and

Metis desired

for coronal

context

radial or

quadrature

alignment

with SPP is

a plus

1.1.2.2 Does

slow and

intermediate

solar wind

originate from

coronal loops

outside of

coronal

holes?

coronal

loops

outside of

coronal

holes

few days

• near

perihelion for

resolution &

better linkage

conditions

• different

phases of

solar cycle

• Mosaic to map

larger region

(e.g. around

AR)

• Modelling to

find best

candidate

source regions

radial

alignment

with SPP is

a plus

1.1.3.2.3 How

does the

heliospheric

magnetic field

disconnect

from the Sun?












5.2.13 L_SMALL_HRES_HCAD_Fast_Wind

Relating observed coronal holes and boundaries to the fast solar wind measured in situ.


Pointing requirements:

Triggers: only IS triggers active (TBC)


PHI PHI_HRT_NOM_0: HRT nominal

mode at 1-min cadence

Flush 1-2 hr period (725MB per hour),

coordinated with EUI.

Note that FDT, highest spatial resolution

(2Kx2K), cadence 6 hours preferred, but

certainly ≥ 1 per day may be useful for context

beforehand (Earth source as well?).

Model PHI flush with 1550MB after SOOP has

completed.

EUI

HRI: 1 – 2 hours, at 1-min cadence,

e.g. EUV & LYA Coronal hole

modes (C) (reduced cadence)

FSI: FSI Synoptic mode (S) (chosen

for model) or FSI Reference Synoptic

mode (R), at several-hour cadence,

both 174 and 304.

≥ 12 hours of much lower-cadence HRI data for

context

We downlink 1-2 hours long period at 1min

cadence (220MB per hour) based on modelling

& LL data from EUI. Coordinated with PHI.

Model EUI flush with 440MB for HRI. FSI

volume based on 2.04kbits/s during whole

SOOP, 3 days =>66Mbytes.

Metis

If polar CH is target, modes are (see

METIS standard modes):

WIND: ≥ 2 hours

MAGTOP: ≥ min. 2 hours; repeated

each day in available observing

windows

FLUCTS: ≥ 1 hour; at perihelion (see

METIS special modes)

Rest of the time Metis will need to

switch off for off-pointing (TBC)

For SAP v0, we can model the Metis

contribution in this SOOP as:

• 2 hours WIND per day

• 2 hours MAGTOP per day

• 1 hour FLUCTS, once in the SOOP (will

be scheduled at closest point)

• Rest of the time Metis off




SPICE

SPICE Composition Mapping: 1

repeat (3.2 hours) (low-latitude CH

only);

SPICE Dynamics: 14 x 4" slit + 8 x

2" slit = 3.72 hours per day

throughout RSW; FOV 8' x 11'

(maybe 4' x 11' at poles)

Possible 4-pixel binning in Y

Default FOV in X for both studies is 4'. Will

assume that these are just placed side-by-side

but with half the time on each position if the CH

is large enough to warrant this. Duration is

therefore unaffected.

Use observation called SPICE_FAST_WIND in

modelling.

SoloHI

Mode: synoptic + turbulence

Currently modelled as

HI_SYN_NEAR +

HI_TURB_PER split to get EID-A

rate

EPD normal+burst

MAG normal+burst

SWA normal+burst

RPW Detection Mode Triggers on



SAP

objective Target

Duration

(milliseconds

)

Opportunity

(e.g., orbital

requirements, solar

cycle phase,

quadrature ...)

Operational constraints Additional

comments



1.1.1.1

Low FIP

fast wind

origins

Sufficientl

y wide

spatial area

across a

CH to

cover

connection

to

spacecraft.

Several hours

of RS

observations

of coronal

source,

considering

time taken for

wind to arrive

at s/c

• Pre-

perihelion

(so, south?)

and

perihelion

parts of RSW

• low-latitude

(to compare

with the intermediate

speed wind observed at

Earth) • any phase of

the cycle, but

more likely

during the

rising phase

• Performed on

several orbits

(preferably;

PHI request

10 orbits)

• EMC quiet for in situ

observations 12 hours

after the RS

observations of the

source region

o can be EMC-

noisy during

RS

observations

Likely to

involve

pointing

away from

disc center

("DC"):

difficult to

see how

Metis can

participate at

those times.

However,

could involve

DC pointing,

especially

between

observations

of the polar

CHs, to look

at

propagating

fluctuations

in/near plane

of sky (but

not

connected to

s/c) at 1.5–

2.9º from

DC.

Doppler

dimming

measurement

s will not be

so useful in these cases,

since the fast

wind will not

have made it

out into the

Metis FOV

(need to check exact

distances)

from the

observed

"source

region".



1.1.1.2

Origin of

the small-

scale X-ray

and UV

jets in polar

coronal

holes

Polar CH

with

sufficiently

wide

extension

to catch

multiple

jets.

as above

• High-latitude

perihelion

(preferred, for fastest

solar wind) • solar

minimum or

declining

phase

as above

Metis will

require

repoints to

DC after off-

center

observations.

1.1.3.2.3

How does

the

heliospheri

c magnetic

field

disconnect

from the

Sun?



5.2.14 L_SMALL_HRES_HCAD_SlowWindConnection

Try to catch with Remote Sensing instruments the dynamics at an open-closed field boundary which

will then be crossed in situ. High resolution RS observations required to catch dynamics. This

SOOP will in general need specific target pointing or target tracking unless the SC is far enough

from the Sun to catch.


Pointing requirements: Off-pointing: specific target pointing or target tracking

Triggers: only IS triggers enabled TBC

See SOOP 2 defined for SOWG8 planning exercise.


EPD close mode +

scheduled/triggered burst

MAG normal +


RPW normal +

scheduled/triggered burst selective downlink useful

SWA normal +




SPICE

combination SPICE

Dynamics & SPICE

Composition Mapping

rasters

• Slit: 4" or 6"

• Exposure time:

180 s


• FoV: (8' / 12') x

11'


6.4 hours

• Nº repeats: 1

Raster area should be optimized to make sure open-closed

field boundary is captured:

• At perihelion if possible, since highest possible

resolution is preferred.

Lines for SPICE Composition Mapping (from 1.1.2.2 Does

slow and intermediate solar wind originate from coronal

loops outside of coronal holes?) :


• Ne VIII 770 Å,

• Ne VIII 780 Å,

• Mg IX 706 Å,

• O II 718 Å,

• O IV 787 Å,

• O V 760 Å,

• O V 761 Å,

• O VI 1032 Å,

• Ne VI 999 Å,

• Ne VI 1010 Å,

• Mg VIII 772 Å,

• Mg VIII 782 Å,

• C III 977 Å,

• Fe III 1017 Å

• Si II 992 Å

N.B. This line list for SPICE Composition Mapping is

similar to that of

L_BOTH_HRES_LCAD_CH_Boundary_Expansion, except

for O VI 1037 (removed) and Si II 992 (added).

Lines for SPICE Dynamics not specified



EUI

EUV & LYA Coronal

hole modes (C) (uses

HRI) at 1min cadence

FSI Synoptic mode (S)

(FSI) throughout

If TM is limited, varying priority schema can be used to

keep HRI data manageable

PHI

Regularly spaced HRT

data at medium to high

resolution

e.g. PHI_HRT_MODE_2

(600s cadence default)

PHI LL magnetograms needed throughout.

Note that 1 hour cadence can suffice for interchange

reconnection at high resolution.

Processing after the RSW: PHI could focus downlink on the

most interesting periods, as inferred from other LL data

Metis

MAGTOP or GLOBAL

(see METIS standard

modes) for context

and linkage of solar wind

source regions to SC

modelled as GLOBAL

with default settings

Only applicable if beyond ~0.5AU during target tracking

SoloHI Context via SoloHI

synoptic modes Model for now as HI_SYN_NEAR

STIX STX_NORMAL

not strictly needed for SOOP although context is

appreciated



SAP


Opportunity

(e.g., orbital

requirements,

solar cycle phase,

quadrature ...)

Operational

constraints

Additiona

l

comments

1.1.4.1.1

Interchange

reconnection

between open

and closed

field lines and

its role in

slow wind

generation

(see Planning

exercise Jan

2016 -

SOOP2)

Open-Closed field line

boundaries (near

ballistic connection

point):

• CH boundaries

• AR edges close

to low-latitude

open field

• Intermediate

areas of quiet

Sun

Target tracking

~1 RSW

(10 days)

• to be

studied for

CHs in

different

locations

(high vs

low

latitudes)

• different

opportuniti

es along

the orbit:

high-

latitude

windows +

perihelion

• to be

studied in

different

solar cycle

phases

• Earth view

before the

observation

s would be

asset to use

modelling

to define

best target

• During

RSW

• pre-

window

synoptics

needed

for target

choice

• VSTP

updates

needed

for target

tracking

• EMC

quiet for

connectiv

ity

radial

alignment

with SPP

is a plus



1.1.2.2 Does

slow and

intermediate

solar wind

originate

from coronal

loops outside

of coronal

holes?

coronal loops outside

of coronal holes few days

• near

perihelion

for

resolution

& better

linkage

conditions

• different

phases of

solar cycle

• it may be

worthwhi

le to map

around

the whole

AR to

have

higher

chance of

being

connected

• EMC

quiet for

connectiv

ity

• Raster

area

should be

optimized

to make

sure

open-

closed

field

boundar

y is

captured

radial

alignment

with SPP

is a plus

1.2.2.6 Study

fast plasma

flows from

the edges of

solar active

regions

discovered

with

Hinode/EIS

edges of solar active

regions - at most likely

ballistic connection

point

few

hours/days

• fast flows

require

high

cadence

observati

ons

(mainly

SPICE

and

HRI?)

radial

alignment

with SPP

is a plus



1.1.2.7 Trace

streamer

blobs and

other

structures

through the

outer corona

and the

heliosphere.

• Near-

quadrature,

so that

SoloHI can

image

Earth-

directed

blobs

objective

also

addressed

by SOOP

y: Link-

outer-

corona-to-

heliospher

e

1.2.2.5

Magnetic

reconnection

in the

chromosphere

, the

transition

region and the

corona

4.4 Are there

separate

dynamo

processes

acting in the

Sun?

in particular:

4.3.1.

Compare the

distribution of

small-scale

fields at low

and high

latitudes

(source:

[SOL-PHI-

MPS-

MN5100-TN-

2])

Quiet Sun at various

latitudes.

several

days

low + high

latitude, to be

repeated along the

cycle

• during

RSW

• feature

tracking

Only PHI,

EUI/HRI

and

SPICE

really

necessary

for this

goal.

PHI/HRT

mode 2

with 5-10

min

cadence, 5

quantities

and no

binning

EUI/HRT

follows

cadence of

PHI

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/SOOP+y%3A+Link-outer-corona-to-heliosphere








1.1.3.2.3 How

does the

heliospheric

magnetic

field

disconnect

from the Sun?

To be discussed whether any of the following can be linked:

1.1.3.1 Full characterization of photospheric magnetic fields and find structures

1.1.3.3 What is the distribution of the open magnetic flux?

1.2.1.7 Detect and characterize waves in closed and open structures



5.2.15 L_BOTH_LRES_MCAD_Pole-to-Pole

This SOOP is designed to be used as a whole- or half-orbit Synoptic campaign that scans the Sun

from one high latitude to the other (therefore mainly to be used later in the NMP at min inclination

of 15º, i.e. orbit 3 or later). Close to the higher latitude windows, we get a radial sweep at nearly

constant latitude for in-situ.

This SOOP resembles very much the L_FULL_LRES_MCAD_Coronal_Synoptic but this time

SPICE is necessary as well.

We could possibly schedule this SOOP after every Venus GAM.


Pointing requirements: mainly disk center



EUI FSI Synoptic mode (S) with FSI 10 min cadence

Metis

One of METIS standard modes to observe

large scale coronal structures: GLOBAL or

LT-CONFIG

+ CMEOBS whenever a CME is triggered

Model: LTCONFIG (cadence 20mins) and 3 x

1 hour CMEOBS

CME Watch On

could do a slight offpoint without

switching metis off, e.g. to poles

during RSW1+3

PHI

FDT in general though HRT could be

scheduled too depending on solar distance, and

e.g. for pole observations

PHI_FDT_MODE_3 / PHI_HRT_MODE_3

(TBC)

at higher latitudes, point to the poles

for polar magnetic field observations.

Can be modelled as

PHI_FDT_MODE_3 everywhere as

resources are same.

SoloHI SoloHI synoptic Operations

Currently modelled as HI_SYN_NEAR



SPICE SPICE Composition Mapping raster followed

by multiple instances of SPICE CME Watch

SPICE scans many latitudes during the

SOOP: large swath around the sun

magnetic field spiral less curved at

higher latitudes:

SPICE naturally closer to connection

point

Use observation called

SPICE_CME_COMP in modelling.

STIX N/A

EPD Normal + Burst Mode

MAG Normal + Burst Mode


SWA Normal + Burst Mode


Opportunity

(e.g., orbital

requirements, solar

cycle phase,

quadrature ...)

Operational

constraints

Additional

comments

2.2.1 How do CMEs

contribute to the global

evolution of magnetic flux

in the heliosphere?

2.2.2 What is the role of

ICMEs in the Sun’s

magnetic cycle?



5.2.16 L_BOTH_MRES_MCAD_Farside_Connection

In this SOOP, the spacecraft points at the modelled ballistic connection point, so this involves

tracking. This is different from slow connection because it is intended to be used when PHI full disk

imagery is needed to see behind the limb as viewed from Earth.


Pointing requirements: Off-pointing


(Will be modelled with duration of 1 day so that it can be repeated as often as needed in a RSW.)


EUI


HRI in EUV & LYA Coronal hole

modes (C)

HRI in CH mode at cadence ~ 800s;

FSI continuously synoptic mode

Metis N/A Door closed

PHI

PHI/FDT synoptic program:

FDT in PHI science mode 6 at low

cadence

at EID-A rate, ideally regular flushes;

Maybe higher spatial resolution than true

synoptics.

SoloHI

SoloHI: Nominal synoptic program

(currently modelled as HI_SYN_NEAR)

SPICE SPICE Composition Mapping & SPICE

Dynamics interleaved


SPICE_WIND_CONNECT in modelling.

(See SPICE Pseudo-observations for SOOPs)


EPD Normal Mode + regular burst mode Close mode until at least the end of RSW

MAG Normal Mode + regular burst mode

RPW Normal Mode + regular burst mode Burst Triggers Active, selective OK.

SWA Normal Mode + regular burst mode





SAP


Opportunity

(e.g., orbital

requirements,

solar cycle

phase,

quadrature

...)

Operational


1.1.2.7 Trace

streamer

blobs and

other

structures

through the

outer corona

and the

heliosphere.

At least

for 8 hrs,

best

several

days.

Near-

quadrature, so

that SoloHI

can image Earth-directed

blobs

objective also addressed by


Add

connectivity

objectives

here



5.2.17 L_BOTH_MRES_MCAD_Flare_SEPs

This SOOP is aimed at understanding SEP properties and dynamics in relation to Flare events. EUI

and STIX are leading this SOOP, while IS payload provides continuous observations. Synoptic

support from other full disk RS instruments. Disk center pointing preferred.

For most of the science objectives SPICE is good to have, but as the SEP events are rather rare, it is

not feasible to go hunting for these events and to point SPICE to any particular region. However,

since all other instruments are involved, SPICE will probably be observing anyway (SPICE best

mode TBC).


Pointing requirements: disk center preferred

Triggers: IS + metis trigger active


EUI (leads)

FSI Synoptic mode (S),

EUV & LYA Active Region modes

(A) (triggered) with 1min cadence

(default = 1or2s)

Most of the objectives in this SOOP are too

exceptional to be hunted for with high-res high-

cadence observations, except when they can be

triggered by a STIX flag whenever the flare is in

the HRI FOV.

Trigger needed: EUI to download only 1 event

in HRI AR mode (~120kbits/s for 1 hour)

STIX

(leads) STIX Normal Mode

Metis

METIS standard

modes GLOBAL,METIS special

modes CMEOBS

CMEOBS starts after CME flag rise, min. obs

time 1 hr (high cadence, 1 min), data volume~

2.137 Gb

SoloHI SoloHI Shock and Synoptic modes

PHI PHI science mode 2 (FDT) FDT in general

SPICE SPICE Waves mode (is sit-and-stare)

SPICE is good to have, but no target hunting

will be performed in this SOOP, so SPICE will

be disk center pointed

MAG MAG Normal Mode

EPD EPD Normal Mode

RPW RPW Normal Mode and RPW Burst

Mode

SWA SWA Normal Mode

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/METIS+standard+modes


https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/METIS+special+modes







SAP


Opportunity

(e.g., orbital

requirements,

solar cycle phase,

quadrature ...)

Operational


3.1.1.6 What

causes SEPs'

spectral

breaks?

Corona +

Heliosphere

Statistics

Radial dependence

Multiple orbits

PHI, STIX and FSI

seem to be the most

valuable RS

instruments to address

this goal.

SoloHI + SPP would

be interesting

combination as well.

EUI/FSI 174, 304, 5

min cadence (FSI

default cadence is 10

minutes)

EUI/HRI 174 & Ly-

alpha if source is

connected to SolO

SPICE would be

helpful if by chance the

SEP source region is in

the HRI FOV AND

EUI/HRI (Lya alpha)

observed it -> not

feasible to go hunting

for these events (TBC

by SWT)







3.1.2.4

Explore the

fact that only

some of the

hard X-ray

peaks are

related to

escaping

electrons,

while others

are not

Corona &

Heliosphere

Statistics

Perihelion Metis

compatible

EUI/HRI 1 min

cadence for 30 min

before and during X-

ray peak (EUI AR

mode triggered based

on STIX flare)

EUI/FSI 10 min

cadence

SoloHI & Metis

connectivity with

Probe.

STIX

RPW

PHI/FDT

3.1.2.5 X-ray

prompt

events

Corona &

Heliosphere

In-situ

observed

electron

spectrum and

hard X-ray

photon

spectrum

correlation

Statistics

Perihelion

TBC: EUI/FDT 10 min

cadence, EUI/HRI 1

min cadence for 30

min before and during

X-ray peak

EUI is required to see

details of the flare

region in order to

decide if electrons and

ions are accelerated by

different processes at

different times (but not

sure how this is done

in practice, we should

discuss what EUI

mode is needed,

synoptic as described

in the objective page

does not seem

adequate).

This question should

be asked to Sam

Krucker.

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/3.1.2.5+X-ray+prompt+events





3.1.2.6

Delayed

events

(between X-

ray peak and

electron

release time)

Corona &

Heliosphere

Explore the

origin of the

delay

It is about a

secondary faint

emission after

the main X-ray

burst that could

be possibly

detected for the

first time.

Statistics

Perihelion

Best when the solar

limb from SolO is

connected to Earth,

or other s/c.

Metis

compatible

TBC: EUI/FDT 2 min

cadence, EUI/HRT 1

min cadence for 30

min before and during

X-ray peak

EUI could be needed

for studying the flare

region. If there is a

CME shock (2nd case)

SoloHI & METIS

could be interesting.

3.1.2.8

Explore the

type III radio

bursts delays

Corona &

Heliosphere

1)Propagation

effects in the

interplanetary

medium.

2)Coronal

magnetic

restructuring in

the aftermath

of CMEs (PHI

needed)

Statistics

Perihelion for

Propagation effects

in the

interplanetary

medium

Track connected

region for many

days at different

distances from Sun

does not seem to be

needed, even the

value of high-res

data for this

objective seems

doubtful (or too

ambitious to plan

for). If it would

drive the SOOP,

this would mean

tracking connected

region for many

orbits, to enhance

the chance of

catching a type III

radio burst.

Metis

compatible

EUI/FDT 174, 304 5

min cadence (EUI

default cadence is 10

minutes)

EUI/HRI 174 & Ly-

alpha 2 min cadence

SPICE again, not to be

hunted for. If SPICE

sees it, that is better but

not driver for the

SOOP.










3.1.3.2

Double-

power law

spectra

Corona +


Metis

compatible

Metis & SoloHI for the

shock observations and

EUI/FSI+STIX for the

flare ones

PHI not needed.

3.2.2

Latitudinal

and

longitudinal

transport of

SEPs

Corona +


Needs range of

latitudes but indeed

this SOOP needs to

be scheduled at

some high-lat

windows as well

Metis

compatible

Best addressed by

multi-viewpoint

statistical dataset, e.g.

SPP (really close to the

sun) + SolO + earth-

based RS data. -> need

many events, viewed

from different

viewpoints and

different distances

Multi-viewpoint

(stereo), thus multiple

SC or Earth assets

SolO: full disk imagery

+ Metis

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/3.2.2+Latitudinal+and+longitudinal+transport+of+SEPs








5.2.18 L_BOTH_HRES_LCAD_CH_Boundary_Expansion

TO BE REVIEWED

SOOP needs review and more elaboration + to be compared to other SOOP

Similar to L_SMALL_HRES_HCAD_SlowWindConnection, but specifically targeted at

overexpanded CH boundaries as slow wind sources, requiring a different PHI mode and SPICE

observations.


Pointing requirements: Off-pointing, combined with disk-center



EPD N/A

MAG normal + scheduled/triggered burst

RPW normal + scheduled/triggered burst

SWA normal + scheduled/triggered burst



SPICE

SPICE Dynamics (on 1 day), 4–6 rasters,

depending on structure. (Modelling as 6

rasters = 1 hour)

3–4 subsequent days of regular

observations: SPICE Composition

Mapping interleaved with SPICE

Dynamics

Full north-south raster only needed for

extended coronal holes. Otherwise point

at part of boundary modelled to be

connection point.

• Slit: 4" or 6"


• X positions: 224 (4") / 160 (6")

• FoV: 16' x 11'

• Observation time: 1.9 / 1.3 hours

per study

• Nº repeats: depends on target

Raster area should be optimized to make sure

open-closed field boundary is captured:

• At perihelion 6 rasters. Full north-

south raster only to be done for

extended holes. Otherwise it is

enough to point at the boundaries, it

could be that only one raster is

enough.

• For stable structures: 3-4 days of

standard observations + 1 day of

mosaic.

• Highest possible resolution is

preferred.

Lines for SPICE Composition Mapping:


• Ne VIII 770 Å,

• Ne VIII 780 Å,

• Mg IX 706 Å,

• O II 718 Å,

• O IV 787 Å,

• O V 760 Å,

• O V 761 Å,

• O VI 1032 Å,

• O VI 1037 Å,

• Ne VI 999 Å,

• Ne VI 1010 Å,

• Mg VIII 772 Å,

• Mg VIII 782 Å,

• C III 977 Å,

• Fe III 1017 Å


– 4 profiles + 6 (2?) intensities:

• H I 1025 Å,

• C III 977 Å,

• O VI 1032 Å,

• Ne VIII 770 Å,

• Mg IX 706 Å,


Estimate for daily mosaic with 4 pointing

positions of full rasters (6" slit, 10s exptime,

few strong lines, 30mins duration): 10 x 4 x

9.84Mb = 0.4Gbits (18% of orbital vol)



EUI

FSI Synoptic mode (S) (FSI174)

Cadence of a few hours around s/c

magnetic connection point

Synoptic observations with FSI174 (perhaps

partial frame)

PHI

Regularly spaced FDT images for

I_cont, B_LOS, gamma and phi.

Cadence of 6 hours

PHI science mode 6 (with 6hrs cadence)

Throughout RSW.

Metis

Global maps of v_outflow in corona

Global n_e maps

Brightness fluctuation spectra

• WIND ≥ 2 hr, at 5 min cadence

• MAGTOP ≥ 2 hrs, twice a day, at

5 min cadence, during perihelion

RSW

• FLUCTS ≥ 1 hr, at 10s cadence,

VL intensity only

SoloHI

Measure solar wind speed above

potential source region in co-ordination

with Metis

STIX N/A






SAP


Opportunity

(e.g., orbital requirements, solar

cycle phase, quadrature ...)

Operational

constraints

Additional

comments

1.1.2.1 Does

slow wind

originate

from the

over-

expanded

edges of

coronal

holes?

Coronal

hole(s),

Coronal

hole

boundary.

few hours

per day

• South RSW

o High latitude

observing is ideal

for polar coronal

holes.

• Perihelion RSW

o good if coronal

hole has equatorial

extension.

• Earth view beneficial

before RS window for

magnetograms to estimate

global solar magnetic field

SPICE

mosaic for

mapping.

Best if s/c is

on same

streamline as

SPP, but not

required



5.2.19 R_FULL_LRES_HCAD_GlobalHelioseismology

From Obj 4 meeting

This SOOP was inspired on the discussion about Global Helioseismology needs during the

Objective 4 meeting in MPS, Göttingen (Oct 2016).

Default SOOP duration: 10 days (TBC)

Pointing requirements: Disk Centre

Triggers: none

SOOP designed for global helioseismology with PHI/FDT only. The PHI configuration is currently

designed to fulfil helioseismology needs for far side imaging - to detect modes passing through the

solar core - and for deep focusing. Note that the specified PHI mode is not (yet) part of the

predefined PHI modes.


PHI

PHI FDT at 1 min cadence, processing

to v_LOS only

(model as PHI_FDT_SYNOPTIC_5

with

parameters CADENCE = 60

[s] COMPR = 8 IC = 0)

Images compressed by binning to 2x2 or

cropping when farther away from the Sun.

Higher compression may possibly be

acceptable for far side imaging.


Opportunity

(e.g., orbital

requirements, solar

cycle phase,

quadrature ...)

Operational

constraints

Additional

comments

4.2.3 Probe the

structure in deep

layers of the Sun -

Deep focusing

Full

Disk

several

days (e.g. 3

days)

Earth (SDO/HMI) -

SC angle between 45

and 60 degrees

10-15Mm

resolution, i.e. 2x2

binning at perihelion

or cropping further

out



4.2.3 Probe the

structure in deep layers

of the Sun - Far side

imaging to detect

modes passing through

the solar core.

Full

Disk

As long as

possible, 60

days would be

ideal but may

not be feasible

Earth (SDO/HMI)

observations in

combination with PHI

observations at far side:

angle 150 to 210

degrees.

Low resolution, v_LOS

only but 1min cadence.

Compression can be

quite high (TBC how

high).



R_FULL_HRES_HCAD_Density_Fluctuations

Study of density fluctuations in the extended corona as a function of the outflow velocity of the

solar wind while evolving in the heliosphere. If SPICE participates, will need to have limb pointing

for a period of time, then return to disc center for the full-Sun instruments.

Default SOOP duration: 8 hours

Pointing requirements: may require offpointing (e.g. limb)

Triggers: only IS triggers enabled (TBC)


EUI


Deep exposures

Long exposures needed to get good SNR

where it overlaps with Metis

Metis

(lead)

FLUCTS for 1 hour:

(20 mins of: 60 x 1s DIT + 10mins

processing x2

then 40mins at 20s cadence)

then MAGTOP (5 to 20 mins cadence)

Extra processing time is only needed at

1s cadence - 60 images can be stored

and queened for processing

primary at perihelion: METIS SNR for

1s cadence + near corotation (8hours)

(latitude changes not a big problem if

~1degree or less)

+ lower cadence during the rest of the

window: several days min to observe the

lower freq

duration: preferred to have it also 8h in

the other 2 RS windows

PHI

PHI science mode 6

6-hr cadence

1024 x 1024

Ic, B, γ, φ

for context



SoloHI

(lead)

Contributes with synoptic+shock+turbulence

modelled as 50% HI_SHOCK_PER + 50%

HI_SYN_NEAR

Seems like this could be a little better

defined.

SPICE

(optional)

SPICE currently proposes to use either

SPICE Limb mode or SPICE Dynamics,

depending on the target.

Limb:

• Slit: 4"



• FoV: 15' x 11'

• Nº repeats: 10


o 18 mins per study

o 3.2 hours total.

Active region (SPICE Dynamics):

• Slit: 4"



• FoV: 8.5' x 11'

• Nº repeats: 1?


• Limb active region target best if

present. Cannot participate if

Metis is observing, so not

modelled for now.

• If at limb: 224 positions of 60-s

exposures, so lasting 4 hours;

otherwise, the duration is as long

as Metis requires.

Lines for SPICE Limb mode:

• H I 1025 Å,

• C III 977 Å,

• O VI 1032 Å,

• Ne VIII 770 Å,

• Mg IX 706 Å,


– 3 profiles and 3 intensities.


• 4 profiles and 6 intensities.


Opportunity

(e.g., orbital

requirements,

solar cycle phase,

quadrature ...)

Operational

constraints

Additional

comments



1.1.2.4 Study of density

fluctuations in the

extended corona as a

function of the outflow

velocity of the solar

wind while evolving in

the heliosphere

Disc-center to

observe whole

corona (except

when SPICE

observes at

limb)

8 hours x

3

1 8-hour

observation in each

RSW of an orbit

Good to have SPP

in quadrature to

observe the

fluctuations in-situ.



R_SMALL_MRES_MCAD_AR_LongTerm

Decay of Active regions: see slides in 4-Bellot-Rubio-SAP4_magnetoconvection5_final.pptx:

Decay process of ARs is not well known:

• Slow, may last a few weeks

• ARs approach limb and suffer from projection effects

• Emerging flux starts to reconnect with preexisting flux very soon

• Appearance of filaments, flux rope eruptions and CMEs in late phases of decay

• Sunspot fragmentation by light bridges?

• Flux erosion by convective flows?

• Role of moving magnetic features?

• How is the AR flux dispersed?

• What is the fate of the flux?


Pointing requirements: target pointing and tracking

Triggers: disabled


EUI EUV & LYA Coronal hole

modes (C) 600 s cadence

EUV + Lyman alpha, FOV and resolution matching

PHI

SPICE SPICE Dynamics (bracketed by

SPICE Composition Mapping) 2" slit, 128 slit positions, scan duration 11 minutes

PHI PHI science mode 2 HRT Full FOV, 10 min cadence, 5 quantities, 2x2 binning

Metis METIS standard modes:

MAGTOP or GLOBAL

Metis can potentially provide context data before

and after the 15 day SOOP (offpointing during the

15 days)




Opportunity

(e.g., orbital

requirements,

solar cycle

phase,

quadrature ...)

Operational


5.5.2.4

Isolated

AR,

complex

AR on E

limb

15 days

Perihelion: to

ensure near-co-

rotation,

Stereoscopy,

which

constrains

Earth-Sun-SC

angle

Duration 15

days

Potentially coordinated with

ground based

DKIST/EST/GREGOR/NST

for short-term studies

1.1.3.3 What is

the distribution

of the open

magnetic flux?

coronal

holes,

QS, AR

1.2.1.3

Contribution of

flare-like events

on all scales

increase cadence; only

partially

1.2.1.7 Detect

and characterize

waves in closed

and open

structures

plumes high latitude

pointing

1.2.2.5 Magnetic

reconnection in

the

chromosphere,

the transition

region and the

corona

increase cadence

(2.1.1.1 CME

initiation)

probably better

FDT



4.1.1.1 Track

granules and

magnetic features

to follow their

motion and

interactions

increase cadence

4.1.3.2 Follow

individual

magnetic features

flux from lower

to high latitudes

4.1.3.4 How

supergranular

flows

facilitate/impede

transport of

magnetic

features?

increase cadence



R_SMALL_HRES_LCAD_Composition_vs_Height

Mapping the abundance of minor ions as a function of height in the corona to distinguish between

slow and fast wind. This will be targeted at the boundary of a streamer, or at an active region on the

limb. SPICE-led.

Default SOOP duration: 3 hrs

Pointing requirements: off-pointing close-in / disk-center farther out

Triggers: disabled


EUI

HRI observations, e.g. in EUV

& LYA Active Region modes

(A). Cadence TBD (currently

modelled with 360s cadence,

i.e. 20.2 kbps)

Context and also higher (than SPICE) cadence

observations in order to interpret the SPICE

composition map

Metis

One of METIS standard modes

before main observations, get

coronal context, but will not

participate at limb pointing

(unless s/c at large heliocentric

distance): mode TBD

Currently modelled as WIND

Context especially important for streamers

PHI

FDT synoptic data:

1024 x 1024 pixels, at 6 hr

cadence

Ic, B, γ, φ

PHI science mode 6

PHI is needed for context magnetic field, but mainly

before or after the SPICE observations as the target

will be on the limb!



SPICE

(lead)

SPICE Composition Mapping

• Slit: 6"

o At larger

distances, when

possible, use

narrow slit (4"),

to enhance

spatial

resolution.


• X positions: 60

• FoV: 6' (or 4') x 11'

• Observation time: 3

hours per repeat

• Nº repeats: depends on

target

Lines for SPICE Composition Mapping:


• Ne VIII 770 Å,

• Ne VIII 780 Å,

• Mg IX 706 Å,

• O II 718 Å,

• O IV 787 Å,

• O V 760 Å,

• O V 761 Å,

• O VI 1032 Å,

• O VI 1037 Å,

• Ne VI 999 Å,

• Ne VI 1010 Å,

• Mg VIII 772 Å,

• Mg VIII 782 Å,

• C III 977 Å,

• Fe III 1017 Å

( L_BOTH_HRES_LCAD_CH_Boundary_Expansion

assumed to be default.)

.


Opportunity

(e.g., orbital

requirements, solar

cycle phase,

quadrature ...)

Operational

constraints

Additional

comments

1.1.2.3

Abundance of

minor ions as a

function of

height in the

corona as

indicator of slow

or fast wind

Limb active

region (incl

the edges),

or streamer

boundary at

limb

3.2 hours,

twice

Any RSW is okay, but

perihelion preferred

for the AR case, and

>0.55AU is preferred

for the streamer case

(requires Metis

compatibility)

put EUI in at least

twice the cadence

of SPICE to

interpret

composition map

(and possible

aliasing in time)





R_SMALL_HRES_LCAD_FineScaleStructure

High resolution observations of ARs or other solar features to discover the finest scales. SOOP with

low cadence observations as we do not aim at analyzing dynamics here. Perihelion preferred.

Pointing requirements: may be run at disk center or offpointed (without metis)

Default SOOP duration: 12 hours (TBC)

Triggers: disabled


EUI

HRI in highest resolution, e.g. in EUV & LYA

Quiet Sun modes (Q) or EUV & LYA Active

Region modes (A) - LOW CADENCE: 10mins

best at perihelion. Model AR mode,

10 mins cadence

PHI

HRT in highest resolution, e.g. PHI science

mode 4 with cadence 10mins

(or PHI science mode 0 with much lower

cadence)

PHI need depending on science goal

SPICE

one of the high res modes, depends on science

goal. Model as SPICE Dynamics (highest

resolution mode).

may be needed for waves and/or

temperature structure discrimination,

depending on science goal

Metis in one of METIS standard modes, e.g.

MAGTOP with 10mins cadence (=default)

may be part of this SOOP for off-limb

observations (close to perihelion

only) like plumes


Opportunity

(e.g., orbital

requirements, solar cycle

phase, quadrature ...)

Operational

constraints

Additional

comments

1.2.1.6 Resolve the

geometry of fine

elemental loop

strands



R_SMALL_HRES_MCAD_PolarObservations

From SOL-PHI-MPS-MN1500-TN-2

This SOOP is inspired by science goal 1 of 2.3 in the MPS document and is also consistent with the

second part of 2.1 (apart from the number of physical quantities)

Default SOOP duration: Several days

Pointing requirements: Poles

Triggers: disabled

SOOP design to address polar magnetic field objectives that don't necessarily rely on the highest

resolution and cadence PHI data, nor all five physical parameters that PHI can return. The rest of the

remote sensing observations provide supporting data.


PHI PHI science mode 2 (HRT) Full FoV, 2-5 minute cadence, 3

quantities, no binning

EUI


EUV & LYA Coronal hole modes (C)

Keep defaults for now, could potentially

match PHI cadence

SPICE SPICE_fast_wind (SPICE Pseudo-

observations for SOOPs)

SoloHI Normal observing programme

Metis Off pointing so door closed

STIX Normal observations


Opportunity

(e.g., orbital

requirements, solar

cycle phase,

quadrature ...)

Operational

constraints

Additional

comments

4.1 How is magnetic

flux transported to

and re-processed at

high solar latitudes?

High

Latitudes



4.2 What are the

properties of the

magnetic field at

high solar latitudes?

Solar

Poles

repeated high

cadence bursts

of several days

duration

high latitudes,

median to high

resolutions

off pointing so

no Metis

4.4 Are there

separate dynamo

processes acting in

the Sun? (4.3.1)

R_SMALL_HRES_HCAD_Photospheric_Dynamics_Structure

SOOP designed to study fine structure in the photosphere, similar to an overall RS 'Burst mode' but

does not involve Metis, SoloHI and STIX.

Default SOOP duration: 1 hr

Offpointing requirements: may be run at disk-center

Triggers: disabled


EUI

EUV & LYA Quiet Sun modes (Q) or EUV & LYA

Coronal hole modes (C)

Model as EUI_HRI_QS with flush volume 2000MB

HRI 1 - 30 s cadence, maybe

interleaved with 0.1s cadence if

TM allows. Lyα

PHI PHI science mode 0 (HRT) 1 min cadence (flush volume

725MB in 1 hr)

SPICE SPICE_WIND_CONNECT (pseudo mode including

SPICE Dynamics and SPICE Composition Mapping)

2" slit, many rasters over similar

FoV to EUI and PHI




Opportunity

(e.g., orbital

requirements,

solar cycle

phase,

quadrature

...)

Operational

constraints

Additional

comments

1.1.3.1 Full characterization of

photospheric magnetic fields and find

structures

Quiet Sun

Coronal

Holes

A few

minutes

to an

hour as often as

we can.

Perihelion for

quiet sun

Close High

Latitudes for

Coronal Holes

Reduced FoV

(1024x1024

for PHI,

equivalent for

others)

1.1.4.1.6 Photospheric reconnection

Quiet Sun

Coronal

Hole

1 hour

per target

Perihelion for

quiet sun

Close High

Latitudes for

Coronal Holes

Full FoV.

Composition

mode may be

useful for

SPICE here

too.

4.2.2 Basic properties of solar high-

latitude magnetic field structures

5.5 Additional Science Objectives of

PHI

5.5.2.5 (Study flux appearance modes

and interactions in QS)

[source: 4-Bellot-Rubio-

SAP4_magnetoconvection5_final.pptx]

Quiet Sun

(disk center,

E/W limb,

N/S polar

region)

3 h per

pointing

(to

secure

good

statistics)

Perihelion for

highest

resolution

First orbit &

later in

mission for

higher

latitudes

Coordination

with DKIST

and/or EST

[alternatively,

GREGOR]




PHI

5.5.2.1 What are the velocity and

magnetic vectors in the solar

photosphere?

[source: SOL-PHI-MPS-MN5100-TN-

2, proposed by Alex Feller]

Disk center

(DKIST/FSP

looking at

solar limb)

few

hours

Perihelion +

quadrature

with Earth

co-

observations

with

DKIST/FSP,

pointing at

limb

Coordination

with DKIST:

FSP

instrument

pointing at

limb

Quadrature

with Earth!

FSP is

designed to

carry out

highly

sensitive

Hanle

diagnostics of

solar limb

magnetic

structures.

These

measurements

can be

calibrated only

if

simultaneous

high resolution

photometric

and Zeeman

measurements

are carried out

from a much

less inclined

vantage point.


PHI

5.5.3.1 Effect of granulation and

oscillations, i.e., interaction of modes

and convection (stereoscopic

helioseismology).

Sunspot few

hours

What is the

max heliocentric

distance that

will give the

required

spatial

resolution?

co-

observations

with NEO at

~30º.

Earth/near-

Earth assets

assumed to be

continuously observing

(HMI, GONG

network, etc.).

If not, then co-

ordination will

need to be

taken into

consideration.



R_SMALL_HRES_HCAD_AR_Dynamics

Tracking a complex AR of the Sun, for AR dynamics or tracking of a region for initiation. Making

use of Solar Orbiter's higher resolution capabilities, so would be at the best resolution (no binnings).

The acceleration mechanism in solar flares, tremendously enhancing (up to factors of ten thousand)

rare elements like 3He and ultra-heavy nuclei, has been puzzling for almost 50 years (e.g., Mason

2007; Reames 2017). The goal of this SOOP is to examine in detail underlying photospheric sources

of these so-called 3He-rich solar energetic particles (SEPs).

Specifically, the SOOP addresses the following two problems (References at bottom of this page):

1) Does the magnetic flux emergence (cancelation) play a fundamental role in energetic particle

production and release from the Sun? What is the growth rate of ARs associated with 3He-rich SEPs

(e.g., rapid growth may imply high intensities/enrichments, shorter time to SEP production/release)?

Statistically, the 3He-rich SEP sources (regions to which we are connected from the Earth) are

located near the west limb (~W55). Due to a projection effects this science question cannot be

properly investigated with NEO based observations. Though STEREO-A was in the right position

having a direct view on 3He-rich SEP sources, it does not provide surface magnetic field data.

Approaching to the Sun the SO connecting point move towards the Sun-SO line, improving

magnetic field observations of the connected regions (Bucik et al. 2014; Chen et al. 2015).

2) Would we detect 3He-rich SEPs from frequent small emerging dipoles (although without

significant EUV flaring as speculated by Wang et al. 2006) at closer distances to Sun? The SO will

be able to detect the events with much smaller intensities than at 1 AU (at perihelion, a factor of ~

50 if applying a simple inverse cube scaling law) probably allowing their detection also during solar

minimum conditions. Thus, with SO we may see 3He-rich SEPs from new sources.

[source: SOL-PHI-MPS-MN5100-TN-2, description updated by Andreas Lagg (email 10 Apr '17)]


Pointing requirements: target tracking

Triggers: enabled


EUI FSI Synoptic mode (S) and EUV & LYA

Active Region modes (A)

Will need triggers to manage TM.

Model flush volume as 2500MB (~1hr

event+FSI)

Metis May contribute if disk-center pointed

(not modelled for now)



PHI

PHI_HRT_NOM_0

+ PHI_HRT_MODE_2 (3 quantities) for

context around 'event'

Model flush volume as 850MB (~1hr

event + rest in mode 2)

SPICE SPICE Composition Mapping & SPICE

CME Watch


SPICE_CME_COMP in modelling.

STIX STIX Normal Mode Triggers Active




Opportunity

(e.g., orbital

requirements,

solar cycle phase,

quadrature ...)

Operational

constraints

Additional

comments

2.1.1.1 CME

initiation

Active Region 2 days

Earth side (≤ 60°),

All phases of the

solar cycle,

perihelion

In RSW

RSW

Extension

needed for

target

selection

VSTP needed

for target

updated

Offpointing so

no Metis

2.1.1.2 CME

structure

Active Region 2 days

Earth side (≤ 60°),

All phases of the

solar cycle,

perihelion

In RSW

RSW

Extension

needed for

target

selection

VSTP needed

for target

updated

Offpointing so

no Metis



3.1.2.1

Understand

energy release

and particle

acceleration

process

Active Region

Properties of

flare energy

release

Acceleration

relating to the

magnetic

reconnection

process

Statistics

(Limited

by EUI

internal

memory)

Perihelion

preferred.

Best when co-

temporal images

from Earth are

available

Target

tracking

EUI/HRI 1 min

cadence over

several hours for

flare with direct

connection to

SolO.

To be done

together with

I_DEFAULT

3.1 How and

where are

energetic

particles

accelerated at

the Sun?

[source: SOL-

PHI-MPS-

MN5100-TN-2]

ARs several

days perihelion

co-

observations

from Earth

5.5 Additional

Science

Objectives of

PHI:

5.5.2.5 How do

magnetic fields

emerge on the

solar surface?

[source: SOL-

PHI-MPS-

MN5100-TN-2]

emerging flux

regions

PHI cadence may

need to be higher

in QS (2-5

minutes) than for

ARs (5-10 mins)

References

Bucik R., Innes D.E., Mall, U. et al. 2014, ApJ 786, 71

Chen N.-H., Bucik R., Innes D.E., Mason, G.M. 2015, A&A 580, A16

Mason G.M. 2007, SSRv 130, 231

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/3.1.2.1+Understand+energy+release+and+particle+acceleration+process








Reames D.V. 2017, Solar Energetic Particles, Lecture Notes in physics 932, Springer

Wang Y.-M., Pick M., Mason G.M. 2006, ApJ 639, 495



R_SMALL_HRES_HCAD_PDF_Mosaic

Looking at the PDF of the magnetic elements. Idea is to scan the solar radius, which is almost 1º,

with a mosaic made up of different positions (3 or 4) from equator to pole.

Inspired on presentation 3-Lagg_polar_magnetic_fields.pdf

Default SOOP duration: ~2 hours

Pointing requirements: mosaic made up of 3 or 4 positions down central meridian

Triggers: disabled

Action on SOC: Compare (and possibly merge) with

R_SMALL_HRES_HCAD_Photospheric_Dynamics_Structure.

(Has similarities with L_SMALL_MRES_MCAD_Connection_Mosaic, too.)


PHI PHI science mode 0 (HRT)

1 minute cadence. Mosaic down the central

meridian, 3 dwell positions, or 4 for some overlap.

At least 10 images per dwell

Overlap in pointings between the measurements to

do.

SPICE

Quick version of SPICE

Composition Mapping (similar to

SPICE mosaic SOOP)

SPICE would require 30 mins at each dwell

position

EUI To be filled in by EUI team - HRI

STIX




Opportunity

(e.g., orbital requirements,

solar cycle phase,

quadrature ...)

Operational

constraints

Additional

comments

4.1.2 Study the subtle

cancellation effects

that lead to the

reversal of the

dominant polarity at

the poles

4.1.4 Study the

influence of

cancellations at all

heights in the

atmosphere

4.2.1 Probability

density function

(PDF) of solar high-

latitude magnetic

field structures.

March or September, to see

the pole that is seen best from

Earth. For the same reason, it

should be able to see the same

target as Earth. High-latitude,

as close as problem.

4.2.2 Basic properties

of solar high-latitude

magnetic field

structures

4.4 Are there separate

dynamo processes

acting in the Sun?



R_SMALL_HRES_HCAD_RSburst

This SOOP describes a coordinated observation of all high resolution RS instruments, running at

highest resolution and variable but high cadence, for a short period of time.

As a planning scenario, we propose to run this SOOP at every perihelion window where we have

some extra TM to spare, or where the campaign would fit without sacrificing too much of the rest of

the orbit.

This SOOP can be run for different targets, also at plain disk center, as it is aimed to discover new

physics and compare high cadence dynamics in all kinds of solar regions.

Default SOOP duration: 10 mins (can be repeated several times when it fits)

Pointing requirements: may be run at disk-center or off-pointed

Triggers: disabled


EUI

HRI high res / high cadence modes,

depending on target:

EUV & LYA Quiet Sun modes (Q),

EUI/HRI Coronal hole mode (C) and

EUI/HRI Active Region mode (A))

EUI/HRI Discovery mode (D) to

"discover" periods <10s, min

observation time 300s.

(currently modelled with AR mode with

default cadence (1&2s), 5400kbps)

QS and AR mode, as defined now, generate at

4600-5400 kbps. (i.e. about 250x EIDA rate)

In most extreme case, i.e. Discovery mode,

HRI would generate 543MB per 10 mins

SOOP duration.

(data rate = 7240 kbps, i.e. about 350x EIDA

rate)

We download all generated data.

PHI PHI science mode 0 for HRT In mode 0, PHI generates at 1607 kbps (about

80x EIDA rate)

SPICE

Depending on science goal, but SPICE

Waves mode for highest resolution sit-

and-stare,

SPICE 30"-wide Movie if you want

more spatial information

In Waves mode, SPICE generates at ~50 kbps

(about 3x EIDA rate). Model as SPICE

Waves mode for SAP v0.









SAP


Opportunity

(e.g., orbital

requirements,

solar cycle

phase,

quadrature

...)

Operational


1.2.1.1 Energy

flux in the lower

atmosphere

Bright source

(could be AR but not

necessary)

20mins to 3/5

hours

perihelion for

highest resolution,

multiple orbits

good to have

Better early

in mission for Ly alpha

degradation.

Duration for EUI observations depend on

mode: discovery mode may not run long but

can be combined with other mode like A or Q.

EUI/HRI-Lya preferred over HRI-EUV.

1.2.1.3

Contribution

of flare-like

events on all

scales

Flaring region

(could be

'quiet' sun for

nanoflares)

2-3 hours near-perihelion Need EPD/SIS as well

1.2.1.4

Observe and

explore flare-

like ‘heating

events’ from

the quiet

corona

Quiet Sun ?

perihelion

preferred but

not required

multiple

orbits good

to have

1.2.1.5

Determine

whether

coronal

heating is

spatially

localized or

uniform, and

time steady or

transient or

impulsive for a wide range

of magnetic

loops with

different

spatial scales.

AR, AR moss,

QS

several

hours

multiple

orbits good

to have

1.2.1.8

Investigate the

role of small

scale magnetic

flux

emergence in

energizing the

above laying

layers

several

hours

near-perihelion

(<0.5AU)

multiple

orbits good

to have

PHI/HRT leading, EUI/HRI required as

support



1.2.2.3 What

are the origins

of waves,

turbulence and

small scale

structures?

to be combined with


(RS cadence TBC!):

• Long-term observations.

• Short duration bursts near perihelion.

1.3.1 Solar

and local

origin of

Alfvénic

fluctuations

spicules above

limb,

AR loops,

CH+boundary

1-2 hours perihelion

Better early

in mission

for Ly alpha

degradation.

to be combined with I_DEFAULT

1.2.1.7 Detect

and

characterize

waves in

closed and

open

structures

spicules at

limb, AR loops



R_SMALL_HRES_HCAD_WaveStereoscopy


Pointing requirements: Off-pointing

Triggers: N/A

The scientific aim is to characterize the properties of waves in the photosphere and their coupling

with the atmosphere.

Waves are one clear mechanism for transferring energy from the photosphere to the chromosphere

and corona. Measuring the properties of the waves requires, in part, a determination of the velocity

field. The line-of-sight velocity component can be determined at different heights in the atmosphere

by observing Doppler shifts in different spectral lines. From the earth’s vantage point we have high

resolution ground based, balloon borne, and satellite instruments. Determining the horizontal

velocity has previously relied on using correlation tracking of intensity variations and rely on the

questionable assumption that the changes in location of the brightness fluctuations reflect the actual

velocity. The orbit and capability to measure Doppler velocities, in conjunction with existing and

upcoming ground-based or near-earth observatories, offers the unique chance to directly measure

two components of the velocity field using the Doppler effect.

High resolution co-temporal measurements including Doppler velocity maps from SO as well as

ground and NEOs are required. In particular, the ground-based and NEOs should include high

resolution Doppler images in the same line (with a higher cadence than that of SO), as well as lines

sampling different heights of the atmosphere. Co-observation with IRIS would be desirable. During

the observing period the earth-Sun-SO angle should be between 30◦ and 60◦ – a range which

represents a compromise between determining the two components of the velocity field and

allowing magnetic features which can act as wave guides to be partially resolved.

For ease of understanding the connection between the different heights, the observations would best

be performed at the center of the disk as observed from the earth (where observations over different

wavelengths are possible). Because also the achievable cadence will be higher on ground than with

SO/PHI, it is preferable to select targets which are closer to the disk center as seen from earth and at

higher heliocentric angles as seen from SO. The highest possible cadence is desirable, and a shorter

time series (of down to 30 minutes of Solar orbiter observations) would still allow the scientific

objectives to be met. (The ground based and NEO should be made for a period of 90 minutes

centered on the 30 minute SO observations). However, in order to guarantee reliable conditions

(seeing) at the coordinating ground-based observing facility (e.g. DKIST) a continuous high-

cadence observation period of several hours is required.

High resolution context magnetic maps from solar-orbiter immediately before and after the 30

minute observing window are required to provide context and aid co-alignment.



A second observational campaign of an area 45◦ from disk center, with an earth-Sun-SO angle of

90◦, would be desirable.


PHI PHI science mode 0 (HRT) full FoV 1 min cadence, 2 quantities plus magnetic

field context

EUI EUV & LYA Quiet Sun

modes (Q)

Full FoV 1-10s cadence (Just Lyman alpha?)


Opportunity

(e.g., orbital

requirements,

solar cycle

phase,

quadrature ...)

Operational

constraints

Additional

comments

1.2.1 What mechanisms

heat the corona?

see

text

above

several

hours, best

30 minutes

downlinked

What is the max

heliocentric

distance that will

give the required

spatial

resolution?

co-

observations

with NEO at

30, 60, 90

degrees.

1.2.1.1 Energy flux in

the lower atmosphere

see

text

above

several

hours, best

30 minutes

downlinked

What is the max

heliocentric

distance that will

give the required

spatial

resolution?

co-

observations

with NEO at

30, 60, 90

degrees.



5.5 Additional Science

Objectives of PHI

5.5.3.2 Two

components of velocity.

What is the relationship

between the components

of the velocities in

granulation?

Supergranulation?

Various modes in quiet

Sun? (stereoscopic

helioseismology)

Quiet

Sun

several

hours

What is the max

heliocentric

distance that will

give the required

spatial

resolution?

co-

observations

with NEO at

~30º.

Earth/near-Earth

assets assumed to

be continuously

observing (HMI,

GONG network,

etc.). If not, then

co-ordination will

need to be taken

into consideration.



R_SMALL_HRES_HCAD_Ephemeral

From SOL-PHI-MPS-MN1500-TN-2

This covers half of SOOP 2.1 from the MPS Document, the other half is covered by

L_SMALL_HRES_HCAD_SlowWindConnection

Default SOOP duration: Several days

Pointing requirements: Quiet Sun

Triggers: disabled

The emergence, diffusion and decay of ephemeral regions near the poles and below high- latitude

coronal holes should be studied for the aspect of how they feed the magnetic network (see e.g.

Simon et al. 2001, ApJ 561, 49 427; Gosic et al. 2014, ApJ 797,). In particular, the latitudinal

dependence of this decay process would be interesting to study.


PHI PHI science mode 3 (HRT) Half FoV 1-2 minute cadence, 5 quantities, no

binning

EUI EUV & LYA Coronal hole modes

(C)

FOV and cadence matching PHI, 15-30s cadence

SPICE SPICE Composition Mapping FOV Matching PHI


Opportunity

(e.g., orbital

requirements, solar

cycle phase,

quadrature ...)

Operational

constraints

Additional

comments

4.1.1 Study the

detailed solar

surface flow

patterns in the polar

regions, including

coronal hole

boundaries.

- - see Helioseismology

SOOPs -



4.4 Are there

separate dynamo

processes acting in

the Sun? (4.3.1.2)

Ephemeral

Regions, (quiet

Sun below the

poles, above

polar coronal

holes).

Several

Days

High latitude, solar

minimum

feature

tracking





6 PLANNING STRATEGY

In order to construct the current version v0 of the SAP, we implemented the planning

strategy described below, by following, step by step, the different presented criteria.

We have to note that the SAP v0 does not include the planning of the Objective 4, even

though all the relevant SOOPs have been defined. However, many of its sub-objectives are already

covered by some of the planned SOOPs. Especially for the strategy of the Objective 4, see 4.0

Overall remarks and feasibility concerning Objective 4 observations with Solar Orbiter. The

remarks from that section will be integrated here at the next version.

Criterion 1: Best resolution RS data at different perihelia through the mission

We schedule an RS burst SOOP (R_SMALL_HRES_HCAD_RSburst) whenever we have a

perihelion with good telemetry downlink. We could aim at different types of targets or even plain

disk center to discover new physics in Solar Orbiter's highest cadence data, also possibly in

unexpected locations.

This campaign would serve several science goals that need very high cadence, need

perihelion and are aiming at different types of regions (see SOOP page). Even if off-pointing is not

possible, this campaign could still be useful to be run on the Sun-disk center region.

The SOOP under consideration is telemetry demanding (see telemetry estimates in SOOP

page R_SMALL_HRES_HCAD_RSburst, 3 to hundreds of times the EID-A rate), so we need good

telemetry at the time of the perihelion or right afterwards.

Alternatively, for some science objectives, e.g. 3.2.6 Effects of energetic particles

propagating downward in the chromosphere, it may be beneficial to schedule a few short RS bursts

into an RS window (or a series of RS windows), to enhance the chances of catching energetic

particles (instead of dedicating all telemetry for high-resolution high-cadence observations during a

few hours of the window).

Implementation of Criterion 1 for SAP v0:1

-> perihelion windows of

• MTP06 - 2021/07/01 - 2022/01/01,

• MTP08 - 2022/07/01 - 2023/01/01,

• MTP10 - 2023/07/01 - 2024/01/01,

• MTP14 - 2025/07/01 - 2026/01/01 (schedule at real perihelion, if you can hold onto the data

until the big underrun)

• MTP20 - 2028/07/01 - 2029/01/01 perihelion though this one if farther out (0.38AU)

1 For the definition and properties of the different MTPs for the October 2018 Option E, see next chapter.



For objectives like 3.2.6 Effects of energetic particles propagating downward in the chromosphere,

we could possibly try to reserve the best telemetry windows for this case:

• MTP08 - 2022/07/01 - 2023/01/01 has 2 concatenated windows at good telemetry, can be

improved even more by shifting both perihelion and north window later, so that perihelion

RSW starts at perihelion.

• Alternatively, MTP14 - 2025/07/01 - 2026/01/01 can be used (similar but at higher

latitudes).

End of the mission may be more restrictive for RS bursts. It is peculiar in the sense that we get

many RSW per MTP but in general somewhat lower telemetry downlink:

- average rate for period 2026-2028 is 1.3 * EIDA rate

- average rate for period 2026-2029 is 1.64 * EIDA rate

Criterion 2: Objectives requiring Metis & SoloHI to observe Earth-directed Transients

The CME structure & propagation objectives as well as the blobs objectives ideally need

Solar Orbiter and Earth in quadrature with SoloHI looking towards Earth, so Solar Orbiter at GSE -

Y.

This criterion can preferably be applied at perihelia but also during high-latitude windows.

Alternatively, instead of quadrature, it will be interesting to observe at 45 degrees separation angle.

SOOPs that are most suitable to run during these times are:

1) L_FULL_HRES_HCAD_Coronal_Dynamics: focused on the off-limb corona up to Earth

2) L_FULL_HRES_HCAD_Eruption_Watch: same as above + the solar disk signatures including

PHI observing at higher cadence to see CME initiation

The second SOOP is more telemetry demanding but helpful if the CME happens to come towards

Solar Orbiter: then it can be viewed sideways from Earth.

Windows that fall close to Equinox could also be preferred since Earth-directed CMEs (and

southward IMF SW-Magnetosphere coupling in general) are more geoeffective (Russell and

McPherron, 1973).

Other SOOPs that should at least be run few times at quadrature

is L_FULL_HRES_MCAD_CME_SEPs and L_FULL_MRES_MCAD_Flare_SEPs (this one with

SoloHI towards Earth). Though these can be run at all times, some of its sub-objectives benefit from

quadrature with Earth, so that Earth can observe the structure of the CME heading towards Solar

Orbiter.

Implementation of Criterion 2 for SAP v0:

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/L_FULL_HRES_MCAD_CME_SEPs?src=contextnavpagetreemode

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/L_FULL_MRES_MCAD_Flare_SEPs?src=contextnavpagetreemode



This happens at the following perihelion windows:

• MTP05 - 2021/01/01 - 2021/07/01

• MTP07 - 2022/01/01 - 2022/07/01

• MTP09 - 2023/01/01 - 2023/07/01 may be less favorable for SoloHI (Thomson sphere) TBC

• MTP13 - 2025/01/01 - 2025/07/01 Perihelion+North window: high telemetry downlink

rate!! + higher latitude at declining phase may be ideal to study polar reversal.

Also, the following RS windows may be interesting, not at quadrature but at 45 degrees separation

angle:

• MTP11 - 2024/01/01 - 2024/07/01 last high latitude window, looking down on North pole at

21º (distance range 0.31-0.34AU)

• MTP 5,7,9 windows happen to fall close to Equinox.

Criterion 3: Slow solar wind connection science requiring Earth context for modelling pre-

RSW

We consider 2 very different types of connection science campaigns each requiring different

contributions from Earth and modelling:

1. During solar minimum, the magnetic field configuration is supposed to be quite simple,

with slow SW coming from the streamer belt. Also during the early orbits of the mission,

Solar Orbiter will stay close to the ecliptic.

If PHI observes the far side magnetic field in good resolution, and we combine that with the

Earth side magnetic field, the full solar magnetic field configuration can be modelled

including the location of the HCS that will determine the hemisphere Solar Orbiter will be

connected to.

This model could be the ideal starting point to do a longer term connection SOOP using

synoptic data of both IS and RS payload pointed to the most likely connection point.

During this campaign, PHI keeps on taking regular full disk magnetograms to update the

magnetic field model as we go. The modelling should also improve as Earth and Solar

Orbiter see overlapping longitude ranges on the Sun.

Proposed planning strategy:

o plan during solar minimum (i.e. early in the mission)

o start with PHI magnetogram data at far side (during one of the higher latitude

windows)

o take some time to construct the model

o use perihelion extension window to update the model and choose the RS target

o keep synoptic program during 10-20 days chasing the connection point

o SOOPs: L_FULL_HRES_LCAD_MagnFieldConfig for the magnetic field modelling

(during first RSW), during the connection observations we use



L_BOTH_MRES_MCAD_Farside_Connection, possibly combined

with L_SMALL_MRES_MCAD_Connection_Mosaic

2. During the rest of the solar cycle, as the Sun becomes more active, also the magnetic field

modelling will become more challenging. In these periods, we hope to rely more on Earth

observations to get a well-constrained model of the field that Solar Orbiter is going to fly

through. If PHI data are restricted or not available, we mainly rely on Earth to produce the

model 4 days in advance due to VSTP turn-around loop. For this to happen, we need Solar

Orbiter in the GSE sector X<1 and Y<0, i.e. similar orbits than the ones needed for Earth-

directed transients above. The further Solar Orbiter moves away from that sector, the more

we rely on PHI data to model the most likely connection point.

Proposed planning strategy:

o start early but also plan later in the mission, towards solar maximum to explore other

types of solar wind source regions

o due to the more complicated and less reliable magnetic field modelling, we may want

to use pointing mosaics to establish the most likely connection

point: L_SMALL_MRES_MCAD_Connection_Mosaic

o during later orbits, the concatenated perihelion+North windows will span a large

range of latitudes over a short period of time (we are currently not sure this is an

asset or rather a complication to connection observation planning)

o SOOPs L_SMALL_MRES_MCAD_Connection_Mosaic combined

with L_SMALL_HRES_HCAD_SlowWindConnection (higher res/cadence RS to

explore source region in more detail) or L_SMALL_MRES_MCAD_Ballistic-

connection

If we point at a coronal hole boundary, also the

SOOP L_BOTH_HRES_LCAD_CH_Boundary_Expansion fits.


Case 1 (solar minimum):

Possible timing: MTP08 - 2022/07/01 - 2023/01/01: PHI observations during South Window

(combined with other science goals), connection science during perihelion+North window based

on magnetic field model.

(Note that MTP06 - 2021/07/01 - 2022/01/01 is a similar orbit but the first RSW needs to be shifted

due to the VGAM. We also do not find more opportunities later in the mission because the Sun will

be more active already and so solar minimum conditions will not be met)

Case 2 (rest of the solar cycle):



Possible timing:

o (MTP05 we skip because we will not offpoint during the first orbit)

o MTP07 - 2022/01/01 - 2022/07/01 perihelion window

o MTP10 - 2023/07/01 - 2024/01/01 2nd South

o MTP11 - 2024/01/01 - 2024/07/01 Last North

o MTP13 - 2025/01/01 - 2025/07/01 (EMP) Perihelion

o MTP15 - 2026/01/01 - 2026/07/01 1st South and 2nd North

o MTP17 - 2027/01/01 - 2027/07/01 could be good if perihelion window gets shifted

earlier a few days: also, good for telemetry

o MTP19 - 2028/01/01 - 2028/07/01 Perihelion

o MTP21 - 2029/01/01 - 2029/07/01 South (+perihelion)

Criterion 4: Polar objectives

The different objectives that require high latitude have to be identified (mainly from Objective 4

that is not included in the current version). They have to be planned during high-latitude windows

and split between objectives that need good telemetry (for a high telemetry window) and those that

they don't.

Partial implementation of Criterion 4 for SAP v0:

MTP20-N + MTP21-S good opportunity for detailed pole analysis, telemetry very good for big

volume of PHI polar data. SOOPs to be added.

Criterion 5: Opportunities for long-term RS observations (concatenated windows or minimal

interruption)

Some science objectives benefit from a longer period of continuous RS observations, typically in

some sort of synoptic mode. A special case of this criterion is RS observations that could run from

pole to pole with minimal interruption.


MTP7 to MTP12 have naturally concatenated RS windows, i.e. 20 days continuous RS window +

possible 4-day extension. These MTPs are favorable

for L_FULL_LRES_MCAD_Coronal_Synoptic SOOP.

Also MTP13 - 2025/01/01 - 2025/07/01 (EMP) could be used as the 4-day extension window

could link 2 RS windows together and give the potential to keep on running (minimal) synoptic

observations.



The following MTP periods are even better in the sense that RS observations can run from pole to

pole with minimal interruption. This is particularly interesting for L_BOTH_LRES_MCAD_Pole-

to-Pole SOOP.

• MTP11 - 2024/01/01 - 2024/07/01 is particularly interesting because it has 2 sets of

concatenated windows. In the 2nd orbit it has only 7 interwindow days between the South

window and the extension start of the concatenated period, so you can observe from 9 May

to 25 June with minimal interruption. You go from -21º to +21º passing through one of the

closest perihelia.

Telemetry is low in that MTP period which is OK since this SOOP has quite moderate

telemetry needs.

• The 2nd orbit in MTP15 - 2026/01/01 - 2026/07/01 (after VGAM) offers the possibility to

observe from 5 May to 19 June with min interruption (few days), flying from pole (31º) to

pole through a perihelion windows of 0.3AU!

• MTP16 - 2026/07/01 - 2027/01/01

• MTP17 - 2027/01/01 - 2027/07/01

• MTP18 - 2027/07/01 - 2028/01/01 - 2nd orbit

Criterion 6: Fast wind connection

L_SMALL_HRES_HCAD_Fast_Wind addresses 2 main science goals that in general need coronal

holes as a target.

Science objective 1.1.1.1 Low FIP fast wind origins would benefit from a low-latitude (or extended)

coronal hole, to increase the chance of connection and to compare the composition of low and fast

solar wind streams: this is most likely to happen in the declining phase of the solar cycle

(Hathaway: DOI 10.1007/lrsp-2015-4).

We prefer orbits/MTP with a fast scan through a big range of latitudes (like the opportunities above

for pole-to-pole SOOP).

In particular, in order to address science goal 1.1.1.2 Origin of the small-scale X-ray and UV jets in

polar coronal holes, high latitude windows are preferred to ensure the presence of a well-established

polar coronal hole that can be observed in full.

Limb pointing from medium latitude is also interesting to get the Doppler velocity component from

SPICE combined with EUI for off-limb intensity. Being up close is a big asset as well!


For fast wind origins (objective 1.1.1.1)

• Best opportunity: MTP15 - 2026/01/01 - 2026/07/01 (2nd orbit, after VGAM) offers the

possibility to observe from 5 May to 19 June with min interruption (few days), flying from

pole (31º) to pole through a perihelion windows of 0.3AU. This MTP has better telemetry

than MTP11, and this is during declining phase. Also, best quadrant for context from Earth.

• MTP13 - 2025/01/01 - 2025/07/01 is also possible, but not very good telemetry



• MTP17 - 2027/01/01 - 2027/07/01

• MTP18 - 2027/07/01 - 2028/01/01 - 1st windows are very good telemetry-wise

(we did not add MTP16 because on far side, and did not add MTP11 because of low telemetry)

For the goal 1.1.1.2:

Closest high-latitude windows are:

• Any North window from MTP11 to MTP18 gives a good range of latitudes and stays within

0.4AU.

• MTP21 - 2029/01/01 - 2029/07/01 - south (up to 31º and 0.38AU)

Criterion 7: Science objectives needing perihelia but low telemetry requirements

L_FULL_MRES_MCAD_Flare_SEPs need medium telemetry downlink. Some of its sub-

objectives require quadrature with Earth, so this SOOP is also mentioned above in Criterion 2.

L_IS_STIX needs low telemetry (in practice this SOOP is likely to run throughout all RS windows).

L_SMALL_HRES_MCAD_Suprathermal_Popul needs perihelion for most of its sub-objectives,

and off-pointing to a target. telemetry needs are low. -> SOOP still needs review and clean-up. Not

yet scheduled in timeline (only 1 of its sub-objectives that need limb pointing).

As the telemetry needs of both SOOPs are moderate to low, we rather schedule at outbound

perihelia.


• MTP11 - 2024/01/01 - 2024/07/01- 1st orbit

• MTP12 - 2024/07/01 - 2025/01/01- End of perihelion with better telemetry rate

• MTP13 - 2025/01/01 - 2025/07/01 (EMP) works as well, also good for quadrature

• MTP15 - 2026/01/01 - 2026/07/01- 1st orbit

• MTP16 - 2026/07/01 - 2027/01/01

• MTP20 - 2028/07/01 - 2029/01/01

Criterion 8: Global magnetic field reconstruction & symmetry

RS windows at the far side of the Sun should be used to have regular, low cadence imaging of

magnetic field, to allow global field reconstruction.

This goal can be addressed by SOOP L_FULL_LRES_MCAD_Coronal_Synoptic

or L_FULL_HRES_LCAD_MagnFieldConfig.

Ideally, we plan this SOOP at regular far side windows covering a wide range of phases in the solar

cycle.

In addition, the same opportunities will satisfy allow L_FULL_LRES_MCAD_Coronal_Synoptic to

address the study of symmetry of the magnetic field and active longitudes (4.2 What are the

properties of the magnetic field at high solar latitudes?).


https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/L_IS_STIX?src=contextnavpagetreemode

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/L_SMALL_HRES_MCAD_Suprathermal_Popul?src=contextnavpagetreemode








At first sight, about 7 windows seem to cover these criteria (examples to be added).

Criterion 9: Rest of the objectives and special circumstances

1.1.2.3 Abundance of minor ions as a function of height in the corona as an indicator of slow

or fast wind

will be addressed through SOOP R_SMALL_HRES_LCAD_Composition_vs_Height:

We need limb pointing to address this science goal: either an AR (with open field at the edges) on

the limb or the boundary of a streamer, so the target can be chosen at the time of VSTP. We prefer

but do not require perihelion. Running this SOOP a bit further out could benefit from Metis

participation and enough signal in SPICE.

Exactly the same requirements are needed for 3.3.1.3 Role of shocks in generating SEPs.

This objective is addressed by SOOP L_SMALL_HRES_MCAD_Suprathermal_Popul and also

needs limb pointing at an RS window above 0.55AU, so that Metis can contribute to the

observations.

Implementation for SAP v0:

There seem to be plenty opportunities to plan this. Currently, this SOOP has been scheduled in

MTPs:

• MTP07 - 2022/01/01 - 2022/07/01

• MTP09 - 2023/01/01 - 2023/07/01

• more to be found

1.2.1.6 Resolve the geometry of fine elemental loop strands

will be addressed by SOOP R_SMALL_HRES_LCAD_FineScaleStructure.

This needs the highest possible resolution at close perihelia, but no high cadence and thus no

particularly high telemetry needs. All close perihelia seem to be candidates for this SOOP.


• MTP11 - 2024/01/01 - 2024/07/01 and MTP12 - 2024/07/01 - 2025/01/01 at 0.28AU

• MTP 7 to 10 and MTP15 to 18 at 0.3AU

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/1.1.2.3+Abundance+of+minor+ions+as+a+function+of+height+in+the+corona+as+indicator+of+slow+or+fast+wind

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/1.1.2.3+Abundance+of+minor+ions+as+a+function+of+height+in+the+corona+as+indicator+of+slow+or+fast+wind


https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/L_SMALL_HRES_MCAD_Suprathermal_Popul?src=contextnavpagetreemode

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/1.2.1.6+Resolve+the+geometry+of+fine+elemental+loop+strands



1.1.2.4 Study of density fluctuations in the extended corona as a function of the outflow

velocity of the solar wind while evolving in the heliosphere

will be addressed by SOOP R_FULL_HRES_HCAD_Density_Fluctuations. Metis and SoloHI are

leading this SOOP.

Telemetry constraints: The SOOP is telemetry limited (less than average) except for Metis that

needs more telemetry and SoloHI that seems to need its average allocation.

This SOOP needs to be repeated at several distances, i.e. each RS window, but not too far out for

Metis to still see the density fluctuations (distance limit to be added! 0.5AU TBC). 8 hour per

window should be enough.


The ideal orbits to tackle this science objective are the ones with 3 RSW that are quite close to the

Sun, e.g. between VGAM 7 and 8 all MTPs have their RS windows within 0.5AU and perihelia at

0.3AU:

• MTP15 - 2026/01/01 - 2026/07/01 (2nd half/orbit is best telemetry-wise)

• MTP16, MTP17, and MTP18

These MTPs all happen to fall in EMP. Another option, in the nominal mission, is to schedule this

SOOP once in every window of

• MTP10 - 2023/07/01 - 2024/01/01 if the option is chosen to move the 3 windows closer

together (mainly south window should be moved closer to make it suitable).

Photospheric dynamics (1.1.3.1 and 1.1.4.1.6)

addressed by SOOP R_SMALL_HRES_HCAD_Photospheric_Dynamics_Structure.

High telemetry needs for EUI, PHI and SPICE during short time (up to 1 hour). We need either

perihelion for quiet Sun or close-in high-latitude windows for coronal holes (i.e. North windows).

For perihelion windows, we can select the same ones as for the RS burst above.


For perihelion windows, we can select the same ones as for the RS burst above (MTP6, 8, 10, 14,

20).

Closest North windows, with reasonable latitude (>20º), are:

• MTP11 - 2024/01/01 - 2024/07/01 (especially start of 2nd north window seems perfect, also

telemetry-wise)

• MTP12 - 2024/07/01 - 2025/01/01 (not good telemetry-wise, so not added to the MTP page

for now)

AR dynamics

SOOP R_SMALL_HRES_HCAD_AR_Dynamics

The best opportunities to study CME initiation and structure (close to the Sun), are to point to ARs

at perihelion. We prefer Earth context for modelling and CME context.






Earth-sided perihelion or close-in windows (which we interpret at <0.4AU) are:

• MTP13 - 2025/01/01 - 2025/07/01 - start of perihelion

• MTP15 - 2026/01/01 - 2026/07/01 - 2nd perihelion & 2nd north RSW

• MTP18 - 2027/07/01 - 2028/01/01 - first North

• MTP19 - 2028/01/01 - 2028/07/01 - perihelion

Opportunities at the edge of 60º angle with Earth are (not added to MTP pages for now):

• MTP07 - 2022/01/01 - 2022/07/01 - start of perihelion

• MTP17 - 2027/01/01 - 2027/07/01 - perihelion

Most of these opportunities are in EMP which seems OK as we will have more chances for CMEs

and dynamic events in general.

3.2.2 Latitudinal and longitudinal transport of SEPs

SOOP L_FULL_MRES_MCAD_Flare_SEPs (important to have SPP data).

It needs many events, viewed from different viewpoints (also separated from Earth) and different

distances. It also needs a range of latitudes (some high-latitude windows as well).

Ideally, it should be scheduled as many periods as possible.


The best opportunities are the following,

• MTP11 - 2024/01/01 - 2024/07/01- Both Perihelia + NW to have different viewpoints with a

range of latitudes and long-term duration (very good telemetry)


range of latitudes and long-term duration (good telemetry)


range of latitudes and long-term duration (bad telemetry in the second Perihelion)

Additionally: MTP14 - 2025/07/01 - 2026/01/01- Perihelion + NW for a range of latitudes and long-

term duration (good telemetry)

Energy flux in the lower atmosphere (1.2.1 What mechanisms heat the corona?)

This requires co-observations with Earth-based (DKIST) and NEO (IRIS) facilities, with sets of

particular geometries between Solar Orbiter, the target on the Sun, and Earth.


https://issues.cosmos.esa.int/solarorbiterwiki/display/~lrodri03/L_FULL_MRES_MCAD_Flare_SEPs







Not added to corresponding MTP pages yet

Subject to the resolution requirements (i.e., minimum distance from the Sun) we suggest the

following opportunities:

• MTP05 - 2021/01/01 - 2021/07/01 - South

• MTP06 - 2021/07/01 - 2022/01/01 - North

• MTP07 - 2022/01/01 - 2022/07/01 - South

• MTP08 - 2022/07/01 - 2023/01/01 - North

• MTP10 - 2023/07/01 - 2024/01/01 - Perihelion + North

• MTP13 - 2025/01/01 - 2025/07/01 (EMP) - South

• MTP14 - 2025/07/01 - 2026/01/01 - North

• MTP15 - 2026/01/01 - 2026/07/01 - Perihelion #2

• MTP18 - 2027/07/01 - 2028/01/01 - Perihelion #1 + North #1

• MTP19 - 2028/01/01 - 2028/07/01 - South

• MTP20 - 2028/07/01 - 2029/01/01 – North

Limb stereoscopy of magnetic fields (5.5.2.1 What are the velocity and magnetic vectors in the

solar photosphere?)

This requires perihelion observations at quadrature so that Earth-based/-orbiting assets (specifically

the DKIST Fast Solar Polarimeter) can measure the magnetic field from an orthogonal

view. R_SMALL_HRES_HCAD_Photospheric_Dynamics_Structure seems like the best fit for this.


Not added to corresponding MTP pages yet

Objectives that could be enhanced with observations from the Parker Solar Probe

• L_FULL_MRES_MCAD_Flare_SEPs. In particular: 3.1.1.6 What causes SEPs' spectral

breaks? and 3.2.2 Latitudinal and longitudinal transport of SEPs

• L_IS_SoloHI_STIX. In particular 3.1.1.2.3 Warped shock fronts and 3.1.1.2.4 Turbulence

and inhomogeneities

• L_FULL_LRES_MCAD_ProbeQuadrature which requires SPP in quadrature with Solar

orbiter

• This list has to be completed.


The associated opportunities have not yet been defined as they need final trajectory of both Solar

Orbiter and Parker Solar Probe. However, the first two SOOPs are scheduled at other opportunities

already (based on the other sub-objectives).


https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/L_IS_SoloHI_STIX?src=contextnavpagetreemode



Objective 4 issues

We have to note that during the SAP v0 the Objective 4 has not been planned, even though many of

its sub-objectives are already covered by some of the planned SOOPs. Especially for the strategy of

the Objective 4, see 4.0 Overall remarks and feasibility concerning Objective 4 observations with

Solar Orbiter. The remarks from that section will be integrated here at the next version

Decay of ARs

•Study long-term behavior of active regions (5.5.2.4)

•Capabilities: near co-rotation, i.e. close to the Sun ideally perihelion and close to ecliptic,

SolO as observatory: we would use SOOP R_SMALL_MRES_MCAD_AR_LongTerm including

PHI, SPICE and EUI,

we also need stereoscopy: combine solar orbiter data with data from ground based observatories

•Target: Isolated active region, complex active region close to E limb


Possible target in option E:

• MTP07 - 2022/01/01 - 2022/07/01 perihelion + RSW3, but better move it earlier (i.e. close

to the Earth) to allow stereoscopy and increase telemetry

• extra repetition in the declining phase: MTP17 - 2027/01/01 - 2027/07/01 first 2

windows, but not in ecliptic (TBC!)

Probability distribution functions of the magnetic elements

These observations require a combination of high solar latitude with low heliocentric distance. In

the example trajectory, this is often in the third (North) RSW of an orbit. Coordinated observations

form Earth-based assets are needed, ideally at a large B0 angle, so close to equinox, particularly

September for the North Solar pole/ March for the South Solar Pole. This is challenging.


This is challenging, although there is one ideal opportunity in Option E.

SOOP R_SMALL_HRES_HCAD_PDF_Mosaic to be run possibly at following opportunities:

• MTP14 - 2025/07/01 - 2026/01/01 - RSW3

Polar observations

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/R_SMALL_MRES_MCAD_AR_LongTerm?src=contextnavpagetreemode

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/R_SMALL_HRES_HCAD_PDF_Mosaic



• R_SMALL_HRES_HCAD_Ephemeral would be run on high-latitude patches of quiet Sun

or underneath high-latitude coronal holes. As this is looking for variations in the lifecycle of

ephemeral regions throughout the cycle (at high latitude), this needs to happen several times

throughout the cycle. This will certainly require off-pointing early in the nominal mission,

and maybe during the extended mission.

In addition, ...

Regular polar observations above 15-20 degrees latitude, ideally covering several phases in the solar

cycle.

Minimal duration of each campaign is 1 day at cadence 60s for v_LOS at highest resolution (in

HRT). Magnetic context once per hour is fine.

Total integrated length should be at least 30 days (more is better, especially for observing cycle

variations).

-> each SOOP using PHI in mode 0 can be used to address these goals (SOOPs to be added)


Not implemented yet.

Deep focusing

To be scheduled at times when Earth (SDO/HMI) - SC angle lies between 45 and 60 degrees.

Run R_FULL_LRES_HCAD_GlobalHelioseismology SOOP during several days (e.g. 3 days).

(FDT at 1 min cadence, only v_LOS. 10-15Mm resolution, i.e. 2x2 binning at perihelion or

cropping further out)

Far side imaging for modes passing through the solar core

Would require Earth (SDO/HMI) observations in combination with PHI observations at far side:

angle 150 to 210 degrees.

Run R_FULL_LRES_HCAD_GlobalHelioseismology as long as possible, 60 days would be ideal

but may not be feasible.

Compression of the images can be quite high (TBC how high).

SOOPs that should be running (quasi-)continuously

• I_DEFAULT -> currently scheduled to run throughout the whole mission. The in-situ

instruments will always contribute to this SOOP + may contribute to other SOOPs at the

same time (all SOOPs starting with 'L_').



• L_IS_STIX -> in SAP v0 we scheduled STIX to run throughout all RS windows in its

default mode. The STIX data volume downloaded can be steered depending on the available

downlink at each time.

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/L_IS_STIX?src=contextnavpagetreemode





7 OCTOBER 2018 OPTION E TRAJECTORY AND MEDIUM TERM PLANNING

7.1 Option E trajectory

Description of the trajectory from CREMA will be written here after the final launch date will be

known. It will subsequently be revised after launch for taking into account the actual orbital

characteristics.

7.2 Planning periods for Option E (MTPs)

The different planning periods of 6 months (called MTPs) for the nominal and extended mission

phases are the following:

NMP planning periods:

• MTP05 - 2021/01/01 - 2021/07/01

• MTP06 - 2021/07/01 - 2022/01/01

• MTP07 - 2022/01/01 - 2022/07/01

• MTP08 - 2022/07/01 - 2023/01/01

• MTP09 - 2023/01/01 - 2023/07/01

• MTP10 - 2023/07/01 - 2024/01/01

• MTP11 - 2024/01/01 - 2024/07/01

• MTP12 - 2024/07/01 - 2025/01/01

EMP planning periods:

• MTP13 - 2025/01/01 - 2025/07/01

• MTP14 - 2025/07/01 - 2026/01/01

• MTP15 - 2026/01/01 - 2026/07/01

• MTP16 - 2026/07/01 - 2027/01/01

• MTP17 - 2027/01/01 - 2027/07/01

• MTP18 - 2027/07/01 - 2028/01/01

• MTP19 - 2028/01/01 - 2028/07/01

• MTP20 - 2028/07/01 - 2029/01/01

• MTP21 - 2029/01/01 - 2029/07/01

In the next sections, we are detailing the characteristics as well as the preliminary science planning

for each one. The MTPs as shown here are based on a pure six-month division (although they may

be adapted where necessary to avoid that RSWs are split across MTPs). However actual MTPs will

follow the ESOC station-scheduling boundaries, which will only approximate this pure six-monthly

division, so some later adaptation of these boundaries is likely.



7.2.1 MTP05 - 2021/01/01 - 2021/07/01

Plots are in GSE (geocentric solar ecliptic) coordinates, so Earth is at [0,0], the Sun is at [1,0]. The plot is the

projection of the orbit on the ecliptic plane.

First orbit in NMP.

7.2.1.1 RS window (default) placement

Period Window/GAM EXT

Start Start End Hcentric Distance Range Hgraphic Latitude Range

SC-

Sun- Earth Angle Conjunction

Safe Mode

Comms

Blackout

Largest

Comms

Roll

GAM

Restrictions

MTP

05

EGAM 1 N/A 2021-

01-01

MTP

05

South 2021-

02-13

2021-

02-17

2021-

02-27 0.71 0.64 0.59 -05.05 -05.12 -05.05 002 001 006 NO NO 00 NO

MTP

05

Perihelion 2021-

03-16

2021-

03-20

2021-

03-30 0.37 0.35 0.36 -00.85 02.12 04.10 059 083 106 NO NO 00 NO

MTP

05

North 2021-

03-30

2021-

04-02

2021-

04-12 0.37 0.43 0.49 04.59 05.12 04.89 119 136 148 NO NO 00 NO

7.2.1.2 Science planning

For the very first NMP orbit, we avoid triggered observations and run more synoptic-like programs.

• I_DEFAULT outside of RS windows, L_IS_STIX on inside all RS windows

• MTP5 - perihelion:

Can address objectives requiring Metis & SoloHI to observe Earth-directed Transients, as

Solar Orbiter is in quadrature with SoloHI looking towards Earth:

https://issues.cosmos.esa.int/solarorbiterwiki/download/attachments/15073452/GSE_orbit_MTP_05.png?version=2&modificationDate=1468247018000&api=v2










o L_FULL_HRES_HCAD_Coronal_Dynamics running for 10 days

(L_FULL_HRES_HCAD_Eruption_Watch would be possible as well, but as it

involves triggers it does not seem ideal for first orbit)

o EUI/HRI and SPICE are not involved in the SOOP above. They can run individual

programs that fit in the TM corridors, and are valuable at disk center. Individual

SOOPs will be assigned later.

Current model: EUI_HRI_QS (EUV_CAD=600 LYA_CAD=600), SPICE_WAVES

Current model (SAP v0):

modelled within the baseline, currently only the perihelion window is filled, as defined above.

+ I_DEFAULT outside of RS windows, L_IS_STIX on inside all RS windows



7.2.2 MTP06 - 2021/07/01 - 2022/01/01

Second NMP orbit, including 3 RS windows.

7.2.2.1 RS windows (default placement):



SC-


Safe Mode

Comms

Blackout

Largest

Comms

Roll

GAM

Restrictions

MTP

06

South 2021-

09-04

2021-

09-08

2021-

09-18 0.70 0.64 0.58 -05.07 -05.12 -05.03 161 164 169 NO

YES (RSW

Days 8 to 10) 00 NO

MTP

06

Perihelion 2021-

10-06

2021-

10-10

2021-

10-20 0.36 0.34 0.37 -00.50 02.47 04.30 134 110 087 NO NO 00 NO

MTP

06

North 2021-

10-20

2021-

10-23

2021-

11-02 0.38 0.44 0.49 04.73 05.11 04.82 075 059 047 NO NO 00 YES

MTP

06

VGAM 4 N/A 2021-

11-23

VGAM4 at 2021-11-23, so 3rd RSW need to moved out of 2021-10-23 to 2021-12-01 period, i.e.

shift 10 days earlier or move to completely different period.

7.2.2.2 New placement for the RS window:

• Keep MTP06 South and Perihelion

• MTP06 North can be moved without losing science as the max inclination of this orbit (<5º)

is negligible anyway. It can be moved either to:











1. After the VGAM, to end 2021 or start 2022, i.e. the very end of MTP06 or start

MTP07. This option has extremely good TM but will be far from the Sun, e.g.

ending around 0.8AU.

2. Move between MTP06 south and perihelion windows. Depending on the science goal

you may want to concatenate it to any of those two, or move all 3 windows together

to cover 30 consecutive days of RS operations (may have drawback for calibration

opportunities though).

• Option 1 has the advantage that it can partially solve the problem of the huge underrun

expected for end 2021-start 2022. Also, the MTP06 perihelion window can be heavily

loaded with TM, stored on-board, and dumped once we have moved closer to Earth.

7.2.2.3 SOOP planning

I_DEFAULT outside of RS windows, L_IS_STIX on inside all RS windows

• MTP06-South: 10 days of Magnetic field synoptics to construct Carrington maps including

Earth data and Solar Orbiter data (still low latitude)

: L_FULL_HRES_LCAD_MagnFieldConfig (addressing objective 1.1.3.2.1 How does the

Sun's magnetic field change over time? & 1.1.3 Source regions of the heliospheric magnetic

field)

• MTP06-Perihelion: Short period (~30mins) of R_SMALL_HRES_HCAD_RSburst

• Alternatively, or combined with previous objective, the South and Perihelion window can

also be merged together and serve to address slow SW connection science goals, in

particular SOOP L_BOTH_MRES_MCAD_Farside_Connection, possibly combined

with L_SMALL_MRES_MCAD_Connection_Mosaic

Note that if you shift the South window later to improve TM and concatenate it with the

Perihelion window, one needs to take into account the conjunction period in between

(concatenated windows may end up being only just 20days long without precursor)

• MTP06-North: to be started asap after GAM, i.e. 30 Nov 2021 running until 9 Dec 2021

(implies NO precursors). Solar distance range of 0.77AU to 0.82AU. To be decided by

SWT whether this is worth pursuing (SoloHI and Metis restrictions).

Current model (SAP v0): RS windows at default locations still

• MTP06_South: L_FULL_HRES_LCAD_MagnFieldConfig throughout RSW

• MTP06_Peri: L_BOTH_MRES_MCAD_Farside_Connection for 9 days & 30

mins R_SMALL_HRES_HCAD_RSburst at end of window

• MTP06-North: L_SMALL_MRES_MCAD_Connection_Mosaic (3hrs/day) combined

with L_SMALL_HRES_HCAD_Fast_Wind rest of the time throughout RSW

(may need to be shortened due to TM restrictions)

• I_DEFAULT outside of RS windows, L_IS_STIX on inside all RS windows










7.2.3 MTP07 - 2022/01/01 - 2022/07/01

3rd orbit in NMP.




SC-


Safe Mode

Comms

Blackout

Largest

Comms

Roll

GAM

Restrictions

MTP

07

South 2022-

02-22

2022-

02-26

2022-

03-08 0.65 0.58 0.52 -12.9 -13.06 -12.77 030 026 020 NO NO 00 NO

MTP

07

Perihelion 2022-

03-22

2022-

03-26

2022-

04-05 0.32 0.30 0.32 -01.81 07.54 12.32 038 069 099 NO NO 00 NO

MTP

07

North N/A 2022-

04-05

2022-

04-15 0.31 0.38 0.44 12.32 12.94 11.53 099 122 138 NO NO 00 NO

7.2.3.2 Current scenario

This orbit has been exercised during the SOWG8 PlanningExercise2016_withResults.pptx.


MTP7 - South:

Start of the RSW is beyond 0.55AU so we can offpoint with Metis still observing. This fits the

requirements of:









• R_SMALL_HRES_LCAD_Composition_vs_Height (streamer case for which we need

Metis)

• L_SMALL_HRES_MCAD_Suprathermal_Popul for objective 3.3.1.3 Role of shocks in

generating SEPs.

We could choose several limb pointings and switch between them during few days.

MTP7 - perihelion:

• R_SMALL_HRES_HCAD_AR_Dynamics would fit at start of perihelion (good TM)

• Objectives requiring Metis & SoloHI to observe Earth-directed Transients, as Solar Orbiter

is in quadrature with SoloHI looking towards Earth:

o L_FULL_HRES_HCAD_Eruption_Watch, or alternatively:

o L_FULL_HRES_HCAD_Coronal_Dynamics (similar as above but less on-disk

activities)

• Combine with connection science goals on slow solar wind sources. These require modelling

to find the most likely connection point.

SOOPs L_SMALL_MRES_MCAD_Connection_Mosaic combined

with L_SMALL_HRES_HCAD_SlowWindConnection (higher res/cadence RS to explore

source region in more detail) or L_SMALL_MRES_MCAD_Ballistic-connection

If we happen to point at a coronal hole boundary, also the

SOOP L_SMALL_HRES_LCAD_CH_Boundary_Expansion fits.

MTP7 - Perihelion+ North: these 2 concatenated windows could also be used

for L_FULL_LRES_MCAD_Coronal_Synoptic SOOP

Modelled in SAP v0:


MTP7 - South: first 2 days of limb pointing:

• 1 day R_SMALL_HRES_LCAD_Composition_vs_Height

• 1 day L_SMALL_HRES_MCAD_Suprathermal_Popul

MTP7 - perihelion: we prefer to model slow wind connection science over eruption watch because of the time in the

solar cycle (rising phase) which simplifies slow wind structure and lower chances for eruptions:

• L_SMALL_MRES_MCAD_Connection_Mosaic (~3hr/day) interleaved

with L_SMALL_HRES_HCAD_SlowWindConnection (~21hrs/day)

(can be modelled in 2 big blocks, i.e. ~1 day of first SOOP and 9 days of 2nd SOOP)

• if it turns out there is extra TM available, we

model R_SMALL_HRES_HCAD_AR_Dynamics on the 1st day of perihelion RSW





https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/L_SMALL_MRES_MCAD_Ballistic-connection

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/L_SMALL_HRES_LCAD_CH_Boundary_Expansion





MTP7 - north: L_FULL_LRES_MCAD_Coronal_Synoptic throughout RSW



7.2.4 MTP08 - 2022/07/01 - 2023/01/01

4rd orbit in NMP.




SC-


Safe Mode

Comms

Blackout

Largest

Comms

Roll

GAM

Restrictions

MTP

08

South 2022-

08-10

2022-

08-14

2022-

08-24 0.65 0.59 0.53 -12.86 -13.07 -12.83 167 171 178

YES (RSW

Days 9 & 10)

YES (RSW

Days 3 to

10)

00 NO

MTP 08

Perihelion 2022-09-06

2022-09-10

2022-09-20

0.32 0.30 0.32 -02.57 06.79 12.03 127 097 066 NO NO 00 NO

MTP

08

North N/A 2022-

09-20

2022-

09-30 0.32 0.37 0.43 12.03 13.00 11.72 066 042 026 NO NO 00 NO

7.2.4.2 Current scenario


Connection science during solar minimum or rising phase, using all 3 RS windows + 1 instance of

the RS burst observations

MTP8 - South window









• PHI magnetogram data at far side using L_FULL_HRES_LCAD_MagnFieldConfig for the

magnetic field modelling

Take some time to combine PHI data with Earth magnetogram data and construct the full solar

magnetic model.

MTP8 - Perihelion

Use perihelion extension window to update the model and choose the RS target:

• keep synoptic program during 10-20 days chasing the connection

point: L_BOTH_MRES_MCAD_Farside_Connection, possibly combined

with L_SMALL_MRES_MCAD_Connection_Mosaic (~3hrs per day)

• At the end of this perihelion window , we also plan a short period (~30mins)

of R_SMALL_HRES_HCAD_RSburst at the chosen target as this perihelion RSW has good

TM downlink which even improves when moving further towards Earth

MTP8 - North

• keep synoptic program chasing the connection

point: L_BOTH_MRES_MCAD_Farside_Connection, possibly combined

with L_SMALL_MRES_MCAD_Connection_Mosaic (~3hrs per day)

Alternative for MTP8 - Perihelion+ North:

• These windows have high TM downlink, and could be suitable to schedule a few instances

of R_SMALL_HRES_HCAD_RSburst to enhance chances of catching energetic particle

events (see 3.2.6 Effects of energetic particles propagating downward in the

chromosphere). TM return can even be improved by moving the 2 windows later,

starting at perihelion instead of before. In the case of starting at Perihelion itself, the

distance range covered during the two concatenated windows would be 0.3-0.5 AU.

• these 2 concatenated windows could also be used


Modelled in SAP v0:

modelled exactly following the above assumptions in first 3 sections (south/peri/north) + I_DEFAULT outside of RS

windows, L_IS_STIX on inside all RS windows






7.2.5 MTP09 - 2023/01/01 - 2023/07/01

5th orbit in NMP, 3 RS windows, reaching max latitude 13º.

RS window (default) placement



SC-


Safe Mode

Comms

Blackout

Largest

Comms

Roll

GAM

Restrictions

MTP 09

South 2023-01-25

2023-01-29

2023-02-08

0.65 0.58 0.52 -12.89 -13.06 -12.77 002 002 008 NO NO 00 NO

MTP

09

Perihelion 2023-

02-22

2023-

02-26

2023-

03-08 0.32 0.30 0.32 -01.88 07.47 12.29 066 097 127 NO NO 00 NO

MTP

09

North N/A 2023-

03-08

2023-

03-18 0.32 0.37 0.44 12.29 12.94 11.55 127 150 166 NO NO 00 NO

Current scenario


• MTP 9 - South: Start of the RSW is beyond 0.55AU so we can offpoint with Metis still

observing. This fits the requirements of:

o R_SMALL_HRES_LCAD_Composition_vs_Height (streamer case for which we

need Metis)

o L_SMALL_HRES_MCAD_Suprathermal_Popul for objective 3.3.1.3 Role of

shocks in generating SEPs.











o L_SMALL_MRES_MCAD_Connection_Mosaic interleaved

with L_SMALL_MRES_MCAD_Ballistic-connection

We could choose several limb pointings and switch between them during few days.

• MTP9 - perihelion: Objectives requiring Metis & SoloHI to observe Earth-directed

Transients, as Solar Orbiter is in quadrature with SoloHI looking towards Earth:

o L_FULL_HRES_HCAD_Eruption_Watch, or alternatively:

o L_FULL_HRES_HCAD_Coronal_Dynamics (similar as above but less on-disk

activities)

• MTP9 - Perihelion+ North: these 2 concatenated windows could also be used

for L_FULL_LRES_MCAD_Coronal_Synoptic SOOP (with current definition, this SOOP

includes L_FULL_HRES_MCAD_CME_SEPs, TBC on SoloHI's contribution)

Modelled in SAP v0:


We are close to solar max so higher chance for eruptions:

MTP9 - South: 2 days of limb pointing (start RSW):

• 1 day R_SMALL_HRES_LCAD_Composition_vs_Height

• 1 day L_SMALL_HRES_MCAD_Suprathermal_Popul

• 1 day L_SMALL_MRES_MCAD_Connection_Mosaic

• 7 days L_SMALL_MRES_MCAD_Ballistic_Connection (in reality the mosaics and ballistic

connection are interleaved, however this is too much detail for modelling at this stage.)

MTP9 - perihelion:

• 1 day of L_FULL_HRES_HCAD_Eruption_Watch at start, and 1 repetition on day 5 -> prioritize 1or2 events in EUI and PHI buffers, to be flushed after 2nd SOOP instance. Metis keeps max 2 CMEs per day.

• Rest of the window (8 days) we schedule L_FULL_LRES_MCAD_Coronal_Synoptic

MTP9 - north: keep L_FULL_LRES_MCAD_Coronal_Synoptic for whole window


https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/L_FULL_HRES_HCAD_Eruption_Watch



7.2.6 MTP10 - 2023/07/01 - 2024/01/01

6th orbit & part of 7th orbit in NMP: 4 RS windows, first S and N windows reach max latitude 13º,

last S window reaches 21º already.




SC-


Safe Mode

Comms

Blackout

Largest

Comms

Roll

GAM

Restrictions

MTP

10

South 2023-

07-13

2023-

07-17

2023

07-27 0.65 0.59 0.53 -12.86 -13.07 -12.83 166 162 155 NO NO 00 NO

MTP 10


2023-08-13

2023-08-23

0.32 0.30 0.32 -02.64 06.71 12.00 100 070 038 NO NO 00 NO

MTP

10

North N/A 2023-

08-23

2023-

09-02 0.32 0.37 0.43 12.00 13.01 11.74 038 014 002 NO NO 00 YES

MTP

10

VGAM 5 N/A 2023-

09-28

MTP

10

South 2023-

12-11

2023-

12-15

2023-

12-25 0.58 0.51 0.45 -20.95 -21.32 -20.47 038 045 055 NO NO 00 NO

MTP10-North placement conflicts with VGAM5

RS window placement

• Move concatenated perihelion and North window earlier by ~5 days to avoid GAM. This

will slightly decrease the downlink rates for these windows, but as we are moving towards

Earth, SSMM can be filled up and all TM can be dumped right after the first 3 RS windows.













• First South window can potentially be moved for better TM if we want to sacrifice the

moderate latitude (13º).

Potentially we could move the 3 RSWs together to make 1 concatenated window : this

could be helpful for L_FULL_LRES_MCAD_Coronal_Synoptic SOOP

and R_FULL_HRES_HCAD_Density_Fluctuations (which has a short duration ~8hrs)

• To check feasibility of this option, we should simulate what happens if

o we give more TM downlink share to IS payload during early MTP09

o let RS produce a lot of TM during MTP09 RS windows (and store on-board)

o lower the TM downlink share to IS payload during the bad comms period

overlapping MTP09 and MTP10 = April-June 2023 (EPD may be in far mode

anyway for last part of that period)

o let all payload produce at regular or better rates during the first 3 MTP10 RS

windows

• If we care more about latitude than solar distance during the 2nd south window, it could

potentially be moved earlier to gain TM. May help moving into MTP11 which has very bad

TM constraints.

Planning


• MTP10-Full First Orbit (S + P + N): move windows together to make 1 concatenated

window and run L_FULL_LRES_MCAD_Coronal_Synoptic SOOP (see above)

o Extra idea for MTP10-Perihelion: Short period (~30mins)

of R_SMALL_HRES_HCAD_RSburst as this perihelion RSW has good TM

downlink which even improves when moving further towards Earth

• MTP10-2nd South: connection science goals on slow solar wind sources. These require

modelling to find the most likely connection point.





SOOP L_SMALL_HRES_LCAD_CH_Boundary_Expansion fits: first time at 20º ->

preferred option for now (first opportunity at 20º

Modelled in SAP v0:


• MTP10-Full First Orbit (S + P +

N): L_FULL_LRES_MCAD_Coronal_Synoptic throughout the 3 RS windows

+ during best TM of perihelion window 10mins of R_SMALL_HRES_HCAD_RSburst










+ on 1st and last day of North window 8 hrs

of R_FULL_HRES_HCAD_Density_Fluctuations

• MTP10-2nd South: L_SMALL_HRES_LCAD_CH_Boundary_Expansion for 10 days




7.2.7 MTP11 - 2024/01/01 - 2024/07/01

7th orbit & most of 8th orbit in NMP: 5 RS windows, S and N windows reach max latitude 21º.




SC-


Safe Mode

Comms

Blackout

Largest

Comms

Roll

GAM

Restrictions

MTP 11


2024-01-07

2024-01-17

0.30 0.28 0.31 -05.99 10.39 19.63 105 137 170 NO NO -01 NO

MTP

11

North N/A 2024-

01-17

2024-

01-27 0.31 0.36 0.43 19.63 21.20 18.96 170 164 148 NO NO 85 NO

MTP

11

South 2024-

05-09

2024-

05-13

2024-

05-23 0.59 0.52 0.46 -20.83 -21.33 -20.71 112 105 094 NO NO 00 NO

MTP

11

Perihelion 2024-

06-01

2024-

06-05

2024-

06-15 0.30 0.28 0.31 -05.50 10.91 19.82 042 010 023 NO NO 00 NO

MTP 11

North N/A 2024-06-15

2024-06-25

0.31 0.37 0.43 19.82 21.16 18.84 023 049 066 NO NO 00 NO


MTP11 - First Perihelion window:

• one of the 3 closest perihelion windows, interesting for high-resolution science

like R_SMALL_HRES_LCAD_FineScaleStructure during few days













https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/R_SMALL_HRES_LCAD_FineScaleStructure



• Metis, SoloHI can benefit from being close in as well:

run L_FULL_LRES_MCAD_Coronal_Synoptic for rest of the time

• (If it fits TM wise, L_FULL_MRES_MCAD_Flare_SEPs could be run here as well

(depending on what has been scheduled in rest of MTP) )

MTP11 - First North window: very close to the Sun at a reasonable latitude

• suitable for L_SMALL_HRES_HCAD_Fast_Wind

MTP11 - Second Perihelion window:

• one of the 3 closest perihelion windows, interesting for high-resolution science

like R_SMALL_HRES_LCAD_FineScaleStructure (can be run on different targets for

example)

• If it fits TM wise, L_FULL_MRES_MCAD_Flare_SEPs could be run here as well

(depending on what has been scheduled in rest of MTP)

• L_BOTH_MRES_MCAD_Flare_SEPs to run through this perihelion and upcoming north

window: benefits of close distance + radial alignment with Earth. SPP would be big asset as

well.

MTP11 - Second North window:

• Objectives requiring Metis & SoloHI to observe Earth-directed Transients, as Solar Orbiter

has 45 degrees separation angle with Earth.

We will be looking down on North pole at 21º (distance range 0.31-0.34AU), around solar

maximum.

SOOP for this RSW:

o L_FULL_HRES_HCAD_Eruption_Watch

o OR L_FULL_HRES_MCAD_CME_SEPs

• Possibly combined with connection science goals on slow solar wind sources. These require

modelling to find the most likely connection point.






• The start of the final north window (or end final perihelion window) is also very close to the

Sun at a reasonable latitude which makes it ideal

for R_SMALL_HRES_HCAD_Photospheric_Dynamics_Structure targeted at a CH. Good

enough TM.

• L_FULL_MRES_MCAD_Flare_SEPs to be continued after perihelion window

• Also, suitable for L_SMALL_HRES_HCAD_Fast_Wind













Alternative plan for full set of second orbit - final 3 RSWs:

• During this period, from 9 May to 25 June, we can run RS observations with minimal

interruption (only 7 interwindow days between the South window and the extension start of

the concatenated period).

SC goes from -21º to +21º passing through one of the closest perihelia. This is particularly

interesting for L_BOTH_LRES_MCAD_Pole-to-Pole SOOP.

TM is low in that MTP period which is OK since this SOOP has quite low TM needs.

Modelled in SAP v0:


MTP11 - First Perihelion window:

• R_SMALL_HRES_LCAD_FineScaleStructure -> duration 2 days, with offpointing

• L_FULL_LRES_MCAD_Coronal_Synoptic -> duration 8 days

MTP11 - First North window: L_SMALL_HRES_HCAD_Fast_Wind duration 3 days - EUI and PHI flush limited

volume (still to be set in SOOP)

MTP11 - 2nd orbit (3 RSWs):

• L_BOTH_LRES_MCAD_Pole-to-Pole SOOP - 30 days in total

• 1 break for R_SMALL_HRES_HCAD_Photospheric_Dynamics_Structure at last day of

perihelion window - duration 1hr





7.2.8 MTP12 - 2024/07/01 - 2025/01/01

end of 8th and 9th orbit in NMP: 3 RS windows, S and N windows reach max latitude 21º, best

perihelia: at 0.28AU.




SC-


Safe Mode

Comms

Blackout

Largest

Comms

Roll

GAM

Restrictions

MTP

12

South 2024-

10-09

2024-

10-09

2024-

10-19 0.59 0.52 0.46 -20.86 -21.33 -20.65 103 110 119 NO NO 00 NO

MTP 12


2024-11-01

2024-11-11

0.31 0.28 0.30 -07.40 08.79 18.97 168 161 127 YES (RSW Days 1 to 4)

YES (RSW Days 1 to 5)

93 NO

MTP

12

North N/A 2024-

11-11

2024-

11-21 0.30 0.36 0.42 18.97 21.30 19.31 127 101 083 NO NO -10 NO

MTP

12

VGAM 6

(EMP) N/A

2024-

12-21


• MTP12 - South window:

• MTP12 - Perihelion window:













o 1 of the 3 closest perihelion windows, interesting for high-resolution science

like R_SMALL_HRES_LCAD_FineScaleStructure

o L_FULL_HRES_LCAD_MagnFieldConfig as we are on the back side: will allow us

to construct 360º magnetic field + very low TM rate, suitable for this orbit

o OR L_FULL_MRES_MCAD_Flare_SEPs (requires medium TM so feasibility to be

modelled)

• MTP12 - North

o suitable for L_SMALL_HRES_HCAD_Fast_Wind

o continue L_FULL_HRES_LCAD_MagnFieldConfig

o could be combined with L_FULL_MRES_MCAD_Flare_SEPs

• MTP12 Perihelion+North: these concatenated windows could be used


Modelled in SAP v0:


south window: /

perihelion: L_FULL_HRES_LCAD_MagnFieldConfig for 10 days

north : L_FULL_HRES_LCAD_MagnFieldConfig (6 days) + L_FULL_MRES_MCAD_Flare_SEPs for

4 days







7.2.9 MTP13 - 2025/01/01 - 2025/07/01 (EMP)

1st orbit in EMP: 3 RS windows, S and N windows reach max latitude 28º.




SC-


Safe Mode

Comms

Blackout

Largest

Comms

Roll

GAM

Restrictions

MTP

13

South 2025-

02-22

2025-

02-26

2025-

03-08 0.59 0.53 0.48 -27.64 -28.17 -27.17 039 032 023 NO NO -08 NO

MTP 13


2025-03-24

2025-04-03

0.34 0.33 0.34 -06.12 10.84 22.44 028 051 075 NO NO 00 NO

MTP

13

North 2025-

04-03

2025-

04-06

2025-

04-16 0.35 0.41 0.46 25.35 28.14 26.16 090 110 126 NO NO 00 NO

• MTP13 Perihelion and North window seem ideal to do CME objectives or any transients

moving towards Earth:

o Solar Orbiter and Earth are in quadrature

o 2 RSW with very good TM downlink









o very high latitude in North window (28º) during solar declining phase: chance of

CMEs affecting polar reversal!

o SOOP: L_FULL_HRES_HCAD_Eruption_Watch

o Also suitable for CME initiation, see

SOOP R_SMALL_HRES_HCAD_AR_Dynamics (close to the Sun at Earth side)

o Alternatively, these 2 concatenated windows could be used


o Also L_FULL_MRES_MCAD_Flare_SEPs would fit as one of its objectives needs

quadrature + L_FULL_MRES_MCAD_CME_SEPs

• MTP13 Perihelion can also be used for connection science goals on slow solar wind

sources. These require modelling to find the most likely connection point.






• MTP13 North also suitable for L_SMALL_HRES_HCAD_Fast_Wind

• MTP13 - full orbit also makes a fast scan through big range of latitudes, and can

alternatively be used for L_SMALL_HRES_HCAD_Fast_Wind to address 1.1.1.1 Low FIP

fast wind origins. But TM is not very good towards the end!

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/R_SMALL_HRES_HCAD_AR_Dynamics







https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/1.1.1.1+Low+FIP+fast+wind+origins

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/1.1.1.1+Low+FIP+fast+wind+origins



7.2.10 MTP14 - 2025/07/01 - 2026/01/01

2nd orbit in EMP: 3 RS windows, S and N windows reach max latitude 28º, farther perihelia: at

0.33AU.




SC-


Safe Mode

Comms

Blackout

Largest

Comms

Roll

GAM

Restrictions

MTP

14

South 2025-

07-22

2025-

07-26

2025-

08-05 0.59 0.53 0.48 -27.67 -28.16 -27.09 176 177 167 NO NO -01 NO

MTP

14

Perihelion 2025-

08-17

2025-

08-21

2025-

08-31 0.34 0.33 0.34 -05.61 11.39 22.78 115 092 068 NO NO 00 NO

MTP

14

North 2025-

08-31

2025-

09-03

2025-

09-13 0.36 0.41 0.46 25.58 28.12 26.04 053 033 018 NO NO 00 NO

• MTP14 - perihelion: another good opportunity to schedule 20-30mins

of R_SMALL_HRES_HCAD_RSburst at perihelion day, at least if you can hold onto the

data until the big underrun. Together with the North window this is a time with very high

TM downlink, and could be suitable to schedule a few instances









of R_SMALL_HRES_HCAD_RSburst to enhance chances of catching energetic particle

events (see 3.2.6 Effects of energetic particles propagating downward in the chromosphere).

• MTP14 - NorthB also suitable for L_SMALL_HRES_HCAD_Fast_Wind (high latitude

and close in)



7.2.11 MTP15 - 2026/01/01 - 2026/07/01

3th and 4th orbit in EMP: 6 RS windows, in the first orbit S and N windows reach max latitude 28º,

in the 2nd orbit they reach 31º already; perihelia at 0.33AU and 0.3AU.




SC-


Safe Mode

Comms

Blackout

Largest

Comms

Roll

GAM

Restrictions

MTP

15

South 2025-

12-19

2025-

12-23

2026-

01-02 0.59 0.54 0.49 -27.51 -28.19 -27.39 028 035 045 NO NO 05 NO

MTP

15

Perihelion 2026-

01-14

2026-

01-18

2026-

01-28 0.34 0.33 0.34 -05.08 11.93 23.10 096 119 143 NO NO 00 NO

MTP

15

North 2026-

01-28

2026-

01-30

2026-

02-09 0.35 0.40 0.45 24.59 28.19 26.51 153 174 170 NO NO 82 NO

MTP

15

VGAM 7 N/A 2026-

03-15

MTP

15

South 2026-

05-05

2026-

05-09

2026-

05-19 0.50 0.43 0.38 -30.01 -31.29 -28.91 113 099 080 NO NO 00 NO

MTP

15

Perihelion 2026-

05-22

2026-

05-26

2026-

06-05 0.31 0.30 0.32 -15.16 06.26 22.12 047 021 006 NO NO 00 NO

MTP

15

North 2026-

06-05

2026-

06-09

2026-

06-19 0.34 0.41 0.46 27.83 31.30 29.59 026 047 061 NO NO 00 NO

















MTP15 - 1st South:

MTP15 - 1st Perihelion:

• L_FULL_MRES_MCAD_Flare_SEPs requires perihelion and medium TM. Close to solar

max so higher chance for flares.

MTP15 - 1st North: also, suitable for L_SMALL_HRES_HCAD_Fast_Wind (high latitude and

close in)

MTP15 - 2nd orbit:

• Use the last orbit (3 windows) for L_BOTH_LRES_MCAD_Pole-to-Pole SOOP.

• Alternatively, or in combination, use those for L_SMALL_HRES_HCAD_Fast_Wind! This

MTP has good TM to support it, it falls during declining phase with higher chance on

extended CHs. Also, best quadrant for context from Earth.

• Alternatively, or in combination, these 3 RSWs can also be favorable to

run R_FULL_HRES_HCAD_Density_Fluctuations with Metis leading (thus disk center

pointing!)

• Perihelion+North is also suitable for L_FULL_MRES_MCAD_Flare_SEPs: benefits of

close distance + radial alignment with Earth. SPP would be big asset as well.

MTP15 - 2nd perihelion & 2nd North:

• close to the Sun at Earth side, so suitable for R_SMALL_HRES_HCAD_AR_Dynamics

MTP15 - 2nd North: also, suitable for L_SMALL_HRES_HCAD_Fast_Wind (high latitude and

close in)





7.2.12 MTP16 - 2026/07/01 - 2027/01/01

5th orbit in EMP: 3 RS windows, S and N windows reach max latitude 31º; perihelia at 0.3AU.




SC-


Safe Mode

Comms

Blackout

Largest

Comms

Roll

GAM

Restrictions

MTP

16

South 2026-

09-16

2026-

09-20

2026-

09-30 0.50 0.43 0.38 -30.08 -31.27 -28.73 116 129 147 NO NO 00 NO

MTP 16


2026-10-07

2026-10-17

0.32 0.30 0.31 -17.54 03.18 20.02 179 156 129 YES (RSW Days 1 to 4)

YES (RSW Days 1 to 4)

-17 NO

MTP

16

North 2026-

10-18

2026-

10-22

2026-

11-01 0.35 0.41 0.47 28.07 31.29 29.49 104 084 069 NO NO 00 NO

MTP16 - full orbit:









• During this orbit, RS observations can run from pole to pole with minimal interruption. This

is particularly interesting for L_Pole-to-Pole SOOP.

• Alternatively, if RS windows can be linked

together: L_FULL_LRES_MCAD_Coronal_Synoptic (also requires long duration RS

observations)

• Other candidate: run R_FULL_HRES_HCAD_Density_Fluctuations at each RS window

(not too far out for Metis to still see the density fluctuations). 8 hour per window should be

enough. Metis is leading, thus disk-center pointing!

MTP16 - South:

MTP16 - Perihelion:

• L_FULL_MRES_MCAD_Flare_SEPs could be run at perihelion. Requires medium TM.

MTP16 - North:

• North window is one of closest high-lat windows, and thus suitable

for L_SMALL_HRES_HCAD_Fast_Wind to address 1.1.1.2 Origin of the small-scale X-ray

and UV jets in polar coronal holes

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/L_Pole-to-Pole

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/L_FULL_LRES_MCAD_Coronal_Synoptic




7.2.13 MTP17 - 2027/01/01 - 2027/07/01

6th orbit + first part of 7th orbit in EMP: 4 RS windows, N and both S windows reach max latitude

31º; perihelion at 0.3AU.




SC-Sun-

Earth

Angle

Conjunction

Safe Mode

Comms

Blackout

Comms

Roll

GAM

Restrictions

MTP

17 South

2027-

01-29

2027-

02-02

2027-

02-12 0.50 0.43 0.37 -30.16 -31.25 -28.54 NO NO NO

MTP

17 Perihelion

2027-

02-15

2027-

02-19

2027-

03-01 0.32 0.30 0.31 -17.03 03.86 20.50 NO NO NO

MTP

17 North

2027-

03-02

2027-

03-06

2027-

03-16 0.35 0.41 0.47 28.30 31.27 29.39 NO NO NO

MTP

17 South

2027-

06-13

2027-

06-17

2027-

06-27 0.50 0.44 0.38 -29.81 -31.31 -29.33 NO NO NO

MTP17-perihelion:



• IF PERIHELION WINDOW GETS SHIFTED FEW DAYS EARLIER (also good for

TM):

connection science goals on slow solar wind sources. These require modelling to find the

most likely connection point.






• close perihelion also candidate for R_SMALL_HRES_LCAD_FineScaleStructure

MTP17 - full orbit (first 3 windows):

• RS observations can run from pole to pole with minimal interruption. This is particularly

interesting for L_BOTH_LRES_MCAD_Pole-to-Pole SOOP.

• Alternatively, or in combination, run L_SMALL_HRES_HCAD_Fast_Wind to address

science goal 1.1.1.1 Low FIP fast wind origins: fast scan through big range of latitudes, and

low-latitude coronal hole would be ideal



enough.

MTP17-North = close high-lat window

• same SOOP as above L_SMALL_HRES_HCAD_Fast_Wind can be used to address science

goal 1.1.1.2 Origin of the small-scale X-ray and UV jets in polar coronal holes







7.2.14 MTP18 - 2027/07/01 - 2028/01/01

second part of 7th orbit + 8th orbit in EMP: 5 RS windows, N and S windows reach max latitude

31º; 2 perihelia at 0.3AU.




SC-Sun-

Earth

Angle

Conjunction

Safe Mode

Comms

Blackout

Comms

Roll

GAM

Restrictions

MTP

18 Perihelion

2027-

06-30

2027-

07-04

2027-

07-14 0.32 0.30 0.32 -16.52 04.54 20.97 NO NO NO

MTP 18

North 2027-07-15

2027-07-19

2027-07-29

0.35 0.41 0.47 28.53 31.25 29.28 NO NO NO

MTP

18 South

2027-

10-26

2027-

10-30

2027-

11-09 0.50 0.44 0.38 -29.89 -31.30 -29.17 NO NO NO

MTP

18 Perihelion

2027-

11-12

2027-

11-16

2027-

11-26 0.32 0.30 0.32 -15.99 05.21 21.43

YES (RSW Days

6 and 7)

YES (RSW

Days 6 and 7) NO

MTP

18 North

2027-

11-26

2027-

11-30

2027-

12-10 0.34 0.41 0.46 27.42 31.31 29.74 NO NO NO

MTP18 - full 8th orbit (i.e. RS windows 3, 4 and 5):



• During this orbit, RS observations can run from pole to pole with minimal interruption. This

is particularly interesting for L_Pole-to-Pole SOOP.

• Alternatively, if RS windows can be linked

together: L_FULL_LRES_MCAD_Coronal_Synoptic (also requires long duration RS

observations)



enough. Metis is leading, thus disk-center pointing!

• Yet other alternative for this orbit with a fast scan through big range of

latitudes: L_SMALL_HRES_HCAD_Fast_Wind to address 1.1.1.1 Low FIP fast wind

origins

MTP18 - First North:

• North windows are one of closest high-lat windows, and thus suitable



• 1st North window also suitable study CME initiation and structure (close to the Sun), with

pointing to ARs. Earth context preferred for modelling and CME

context: R_SMALL_HRES_HCAD_AR_Dynamics

• L_FULL_MRES_MCAD_Flare_SEPs : benefits of close distance + high latitude + radial

alignment with Earth. SPP would be big asset as well.

MTP18 - Second North:

• As above, also this North window is one of closest high-lat windows, and thus suitable



https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/L_Pole-to-Pole

https://issues.cosmos.esa.int/solarorbiterwiki/display/SOSP/L_FULL_LRES_MCAD_Coronal_Synoptic




7.2.15 MTP19 - 2028/01/01 - 2028/07/01

9th orbit in EMP: 3 RS windows, N and S windows reach max latitude ~33º; relatively far

perihelion at 0.38AU.


MTP19-Perihelion:

• connection science goals on slow solar wind sources. These require modelling to find the













• perihelion window also suitable to study CME initiation and structure (close to the Sun),

with pointing to ARs. Earth context preferred for modelling and CME

context: R_SMALL_HRES_HCAD_AR_Dynamics



7.2.16 MTP20 - 2028/07/01 - 2029/01/01

10th orbit in EMP: 3 RS windows, N and S windows reach max latitude ~33º; relatively far

perihelion at 0.38AU.


MTP20-perihelion:

• RS burst SOOP (R_SMALL_HRES_HCAD_RSburst). We could aim at different types of targets or even plain disk center to

discover new physics in Solar orbiter's highest cadence data, also possibly in unexpected locations. Note that this perihelion is relatively far out: at 0.38AU.

• L_FULL_MRES_MCAD_Flare_SEPs

MTP20-North:




• good opportunity for detailed pole analysis, TM very good for big volume of PHI polar

data. SOOP to be added

7.2.17 MTP21 - 2029/01/01 - 2029/07/01

11th and last orbit in EMP: 3 RS windows, N and S windows reach max latitude ~33º; relatively

far perihelion at 0.38AU.


MTP21-South(+Perihelion):



• connection science goals on slow solar wind sources. These require modelling to find the







Could be extended to perihelion window

• South window also good opportunity for detailed pole analysis, TM very good for big

volume of PHI polar data. SOOP to be added

• Close high-lat window, interesting for L_SMALL_HRES_HCAD_Fast_Wind, to address

science goal 1.1.1.2 Origin of the small-scale X-ray and UV jets in polar coronal holes

MTP21-Perihelion:

• see above: connection science goals on slow solar wind sources. These require modelling to

find the most likely connection point.
















8 SIMULATIONS

This section will contain all simulations performed once the full mission timeline is known. This

will be done after the knowledge of the launch date and final orbit is stabilized. The simulations will

include the SSMM fill state, the fill state per instrument stores, power consumption etc.



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