Solar/Stellar Winds: An Overview Steven R. Cranmer Harvard-Smithsonian Center for Astrophysics (Untapped gold mines for space physicists?)

Solar/Stellar Winds: An Overview

Steven R. CranmerHarvard-Smithsonian Center for Astrophysics

(Untapped gold mines for space physicists?)

Solar/Stellar Winds: An Overview S. R. Cranmer, NESSCOctober 27, 2008

Outline

The Stellar Zoo:

• The Sun

• Young low-mass stars (T Tauri, FU Ori, etc.)

• Evolved low-mass stars (giants, supergiants, Miras)

• High-mass stars (Wolf-Rayet, O, B, A, LBVs)

Universal Physical Processes:

• Pulsations / waves / shocks / turbulence

• Magnetic fields: dynamos, reconnection between open & closed regions

• “Radiation hydrodynamics:” forcing, heating/cooling, kinetic effects?


Stellar winds on the HR Diagram

30,000 10,000 6,000 3,000

no coronae?

"hot" solar-type

winds

"warm" hybrid winds

"cool" dense

(slow?) winds

Sun

radiatively driven winds

Be stars

I

III

V

flare stars

106

104

102

1

10–2

O B A F G K M


Driving a stellar wind

• Gravity must be counteracted above the photosphere (not below) by some continuously operating outward force . . .

Gas pressure: needs T ~ 106 K (“coronal heating”)

Radiation pressure: possibly important when L* > 100 L

Wave pressure: can produce outward acceleration in a time-averaged sense

Magnetic buoyancy: plasmoids can be “pinched” like melon seeds and carry along some of the surrounding material . . .

• ion opacity? (Teff > 15,000 K)

• free electron (Thomson) opacity? (goes as 1/r2 ; needs to be supplemented)

• dust opacity? (Teff < 3,500 K)

~

~

F = ma


The solar wind: very brief history• Mariner 2 (1962): first direct confirmation of continuous supersonic solar wind,

validating Parker’s (1958) model of a gas-pressure driven wind.

• Helios probed in to 0.3 AU, Voyager continues past 100+ AU.

• Ulysses (1990s) left the ecliptic; provided 3D view of the wind’s connection to the Sun’s magnetic geometry.

• SOHO gave us new views of “source regions” of solar wind and the physical processes that accelerate it . . .


The coronal heating problem• We still don’t understand the physical processes responsible for heating up the

coronal plasma. A lot (not all!) of the heating occurs in a narrow “shell.”

• Most suggested ideas involve 3 general steps:

1. Churning convective motions that tangle up magnetic fields on the surface.

2. Energy is “stored” (above the photosphere) in the magnetic field.

3. Energy is released as heat, either via particle-particle or wave-particle “collisions.”

Heating Solar wind acceleration


The solar wind acceleration debate

vs.

• What determines how much energy and momentum goes into the solar wind?

Waves & turbulence input from below?

Reconnection & mass input from loops?


The solar wind acceleration debate

vs.

• What determines how much energy and momentum goes into the solar wind?

Waves & turbulence input from below?

Reconnection & mass input from loops?

• Cranmer et al. (2007) explored the wave/turbulence paradigm with self-consistent 1D models of individual open flux tubes.

• Boundary conditions imposed only at the photosphere (no arbitrary “heating functions”).

• Wind acceleration determined by a combination of magnetic flux-tube geometry, gradual Alfvén-wave reflection, and outward wave pressure.



30,000 10,000 6,000 3,000

no coronae?

"hot" solar-type

winds

"warm" hybrid winds

"cool" dense

(slow?) winds

Sun


Be stars

I

III

V

flare stars

106

104

102

1

10–2

O B A F G K M


Pre-Main-Sequence accretion phases

Feigelson & Montmerle

(1999)

M. Burton (UNSW)


T Tauri stars: accreting young Suns

(Matt & Pudritz 2005, 2008)

(Romanova et al. 2007)

• T Tauri stars exhibit signatures of disk accretion (outer parts), “magnetospheric accretion streams” & X-ray corona (inner parts), and various (polar?) outflows.

• Nearly every observational diagnostic varies in time, sometimes with stellar rotation, but often more irregularly.


Accretion-driven T Tauri winds• Recent work has extended the solar wave/turbulence models to outer atmospheres

of young, solar-type stars (Cranmer 2008, arXiv:0808.2250).

• The impact of inhomogeneous “clumps” on the stellar surface generates MHD waves that propagate horizontally (like solar Moreton/EIT waves?).

• These “extra” waves input orders of magnitude more energy into a turbulent MHD cascade, and can give rise to stellar winds with dM/dt up to 106 times solar!



30,000 10,000 6,000 3,000

no coronae?

"hot" solar-type

winds

"warm" hybrid winds

"cool" dense

(slow?) winds

Sun


Be stars

I

III

V

flare stars

106

104

102

1

10–2

O B A F G K M


Cool-star mass loss rates


Cool-star dimensional analysis . . .

• Stellar wind power:

• Schröder & Cuntz (2005) investigated an explanation via convective turbulence generating atmospheric waves . . .

• Funny things happen during rapid evolutionary stages! (e.g., Willson 2000, Ann. Rev.)

• Reimers (1975, 1977) proposed a semi-empirical scaling:


Schröder & Cuntz (2005)scaling for lum. classes I, III, V

Cool-star mass loss rates


How can massive winds be “cold?”• The extended solar corona is so low-density, the conservation of internal energy is

essentially a balance between local heating, downward conduction, and upward adiabatic losses.

• When the outer atmosphere becomes massive enough, though, radiative cooling [~ρ2 Λ(T)] becomes more efficient throughout the wind:

• The high-density wind becomes an extended chromosphere (supported by wave pressure??).

• For this case, Holzer et al. (1983) showed the energy equation is ~irrelevant in determining mass flux! A simple analytic model (of the momentum equation) suffices.


Cool stars: Mira supergiants• How are the presumably cool (“corona-free”)

winds of red giants and supergiants actually accelerated?

• How do these winds affect the shapes of the planetary nebulae that are formed at the end of stellar evolution?

• High-luminosity: radiative driving... of dust?

• Shock-heated “calorispheres?” (Willson 2000)



30,000 10,000 6,000 3,000

no coronae?

"hot" solar-type

winds

"warm" hybrid winds

"cool" dense

(slow?) winds

Sun


Be stars

I

III

V

flare stars

106

104

102

1

10–2

O B A F G K M


Hot star winds: radiative driving

• Castor, Abbott, & Klein (1975) worked out how a hot star’s radiation field can accelerate a time-steady wind, even if its “Eddington factor” is << 1.

• Bound electron resonances have higher cross-sections than free electrons (i.e., spectral lines dominate the opacity)

• In the accelerating wind, these narrow opacity sources become Doppler shifted with respect to the star’s photospheric spectrum.

Kudritzki et al.

• Acceleration depends on velocity & velocity gradient! (Turns F=ma on its head? Nonlinear feedback...)

• Wind momentum (~mass loss rate times terminal speed) scales as a “universal” power law function of stellar luminosity.


Hot star winds: pulsations & waves• O, B, A type stars have small convective cores, but not a

subsurface convection zone (i.e., no corona).

They exhibit “g-mode” pulsations, with velocity amplitudes often exceeding the sound speed at the photosphere!

Can the pulsational energy “leak” out of the star into the stellar wind? It seems to be doing it in at least one case (BW Vul); probably many others.

Some stars rotate so rapidly they become oblate, with some models predicting convection and turbulence at equators.

Massa 1994

MPG courtesy John Telting?


Hot star winds: magnetic channeling?

• O,B-star winds are highly supersonic (v > 1000 km/s) from very close to the stellar surface.

• Ud-Doula & Owocki (2002, 2008) modeled the interaction between this kind of wind and a dipole surface field.

fast

er r

otat

ion

stronger field

• Collisions of parcels lead to stochastic upflows/downflows.

• Fast rotation: a rigid “magnetosphere” develops (seen in Ap/Bp stars?)


• “You do this with telescopes?!” (observational diagnostics of stellar mass loss)

• When is a temperature not the temperature? (preferential ion heating & kinetic effects)

• Accretion disks as torque wrenches (T Tauri)

• Excretion disks as torque wrenches!! (Be stars)

click

click

click

click

Topics for discussion (later?):

Conclusions

• Positive feedback and collaboration between the solar physics, space physics, plasma physics, and astrophysics communities.

• What do you find interesting?

A goal of this round of NESSC:


Extra slides . . .


Diagnostics of stellar mass loss

(Bernat 1976)

• Optical/UV spectroscopy: simple blueshifts or full “P Cygni” profiles

• IR continuum: circumstellar dust causes SED excess

• Molecular lines (mm, sub-mm): CO, OH maser

(van den Oord & Doyle 1997)

wind

star

• Radio: free-free emission from (partially ionized?) components of the wind

• Continuum methods need V from another diagnostic to get mass loss rate.

•

• Clumping?


Multi-line spectroscopy• 1990s: more self-consistent treatments of radiative transfer AND better data

(GHRS, FUSE, high-spectral-res ground-based) led to better stellar wind diagnostic techniques!

• A nice example: He I 10830 Å for TW Hya (pole-on T Tauri star) . . .

Dupree et al.

(2006)


The solar wind mass loss rate• The sphere-averaged “M” isn’t usually considered by solar physicists.

• Wang (1998, CS10) used empirical relationships between B-field, wind speed, and density to reconstruct M over two solar cycles.

ACE (in ecliptic)


Sun’s mass loss history• Did liquid water exist on Earth 4 Gyr ago? If “standard” solar models are correct,

a strong greenhouse effect was needed.

• Sackmann & Boothroyd (2003) argued that a more massive (~1.07 M) young Sun could have been luminous enough to solve this problem, but it would have needed strong early mass loss . . .

Sackmann & Boothroyd (2003)

M ~ LX0.4

M ~ LX1.3

M ~ LX0.1

M ~ LX1.0


Multi-fluid collisionless effects?

Polar coronal hole model


Multi-fluid collisionless effects?

protons

electrons(thermal core only)

O+5

O+6


Departures from thermal equilibrium• UVCS/SOHO observations rekindled theoretical efforts to understand collisionless

heating and acceleration effects in the extended corona.

Alfven wave’s oscillating

E and B fields

ion’s Larmor motion around radial B-field

• Ion cyclotron waves (10–10,000 Hz) suggested as a “natural” energy source that can be tapped to preferentially heat & accelerate heavy ions.

MHD turbulencecyclotron resonance-

like phenomenasomething

else?


Angular momentum removal

• Many accreting T Tauri stars are slowly rotating, despite the fact that disk accretion adds angular momentum to the star (e.g., Bouvier et al. 1993; Edwards et al. 1993).

• How is angular momentum carried away?

CTTS

(accreting)

WTTS

(non-accreting)

Rot. Period (days)

By field lines that thread the disk? (“disk locking”) This would imply that magnetic reconnections (and X-rays!) scale with accretion rate. No...

By CME-like ejections from the tangled field in the disk?

By a stellar wind! Matt & Pudritz (2005, 2008) say Mwind ~ 0.1 Macc can do the job.

L. Hartmann, lecture notes


What kind of outflow is it really?

• YSOs (Class I & II) show jets that remain collimated far away (AU → pc!) from the central star.

• However, EUV emission lines and He I 10830 P Cygni profiles indicate the blueshifted outflow is close to the star.

• Stellar winds & disk winds may co-exist.

(Ferreira et al. 2006)


Be stars: “decretion disks”

• Be stars are non-supergiant B-type stars with emission in hydrogen Balmer lines.

• Be stars are rapid rotators, but are not rotating at “critical” / “breakup:”

Vrot (0.5 to 0.9) Vcrit

• How does angular momentum get added to the circumstellar gas ?

Hints:

• Many (all?) Be stars undergo nonradial pulsations (NRPs).

• Rivinius et al. (1998, 2001) found correlations between emission-line “outbursts” and constructive interference (“beating”) between NRP periods.

• Ando (1986) & Saio (1994) suggested that some NRPs could transfer angular momentum outwards: similar to wave “radiation pressure” (Cranmer, in prep.)

Documents

Solar/Stellar Winds: An Overview Steven R. Cranmer Harvard-Smithsonian Center for Astrophysics (Untapped gold mines for space physicists?)