Leaky Stars: Steven R. Cranmer & many others Harvard-Smithsonian Center for Astrophysics Pulsations, Waves, and Turbulence in Stellar Winds across the

Leaky Stars:

Steven R. Cranmer & many othersHarvard-Smithsonian Center for Astrophysics

Pulsations, Waves, and Turbulencein Stellar Winds

across the H-R Diagram

Leaky Stars:

Steven R. Cranmer & many othersHarvard-Smithsonian Center for Astrophysics

Pulsations, Waves, and Turbulencein Stellar Winds

across the H-R Diagram

Outline:

• Background: history & basic physics

• The Sun: coronal heating & fast solar wind

• Hot stars (O, B, W-R): pulsations & radiation-driving

• Cool stars (T Tau, Mira): chromospheric flows?

Leaky Stars: Pulsations, Waves, and Turbulence in Stellar Winds across the H-R Diagram

S. R. CranmerCfA Colloquium, March 9, 2006

Motivations . . .

• “Space weather” can affect satellites, power grids, and astronaut safety.

• Sun’s mass-loss history may have impacted planetary formation / atmospheres.

• The Sun is a “benchmark” for many basic processes in plasma physics.

Solar corona & wind:

• Mass loss affects evolutionary tracks (isochrones, cluster HB/RGB), SN yields.

• Hot-star winds influence ISM abundances & ionization state of Galaxy.

• Spectroscopy of wind lines extragalactic standard candles?

Stellar winds:



• Coronae & Aurorae seen since antiquity . . .

First observations of stellar outflows ?

• “New stars”

1572: Tycho’s supernova

1600: P Cygni outburst (“Revenante of the Swan”)

1604: Kepler’s supernova in “Serepentarius”



Brief history: solar wind• 1860–1950: Evidence slowly builds for outflowing magnetized plasma in the

solar system:

• 1958: Eugene Parker proposed that the hot corona provides enough gas pressure to counteract gravity and accelerate a “solar wind.”

• 1962: Mariner 2 provided direct confirmation!

• solar flares aurora, telegraph snafus, geomagnetic “storms”• comet ion tails point anti-sunward (no matter comet’s motion)



Brief history: stellar winds

O supergiant (Morton 1967)

M supergiant (Bernat 1976)

• Milne (1924): rad. pressure can eject atoms/ions from stellar atmospheres.

• P Cygni profiles = winds:

» O, B, WR, LBVs: Beals (1929); Swings & Struve (1940)

» G, K, M giants, supergiants: Adams & MacCormack (1935); Deutsch (1956)

• Also: IR excesses, maser emission, “plain” blueshifts.



Schematic H-R Diagram

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

Sun

I

III

V

106

104

102

1

10–2 O B A F G K M



Stellar winds

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



Convection zones

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

Sun

I

III

V

106

104

102

1

10–2 O B A F G K M

deepcore

convection

subsurfaceconvection

fullyconvective



Temperatures in outer atmospheres

Sun

Hot star (O, B)

Cool star (K, M)

?



One-page stellar wind physics

• Momentum conservation:

To sustain a wind, /t = 0 , and RHS must be “tuned:”





• Energy conservation:








• Photosphere (& most of hot-star wind)



• Chromosphere: heating rad. losses














• Transition region & low corona









• Transition region & low corona

• Extended corona & cool-star wind



Solar convection & surface waves• Cool stars with sub-photospheric convection undergo “p-mode” oscillations:

• Lighthill (1952) showed how turbulent motions generate acoustic power; more recently generalized to MHD . . .



Solar convection & surface waves• Cool stars with sub-photospheric convection undergo “p-mode” oscillations:

• Lighthill (1952) showed how turbulent motions generate acoustic power; more recently generalized to MHD . . .



Coronal heating mechanisms• A surplus of proposed models! (Mandrini et al. 2000; Aschwanden et al. 2001)

• Where does the mechanical energy come from?

• How is this energy coupled to the coronal plasma?

• How is the energy dissipated and converted to heat?

wavesshockseddies

(“AC”)

vs.

twistingbraiding

shear

(“DC”)vs.

reconnectionturbulenceinteract with

inhomog./nonlin.

collisions (visc, cond, resist, friction) or collisionless



Coronal heating mechanisms• A surplus of proposed models! (Mandrini et al. 2000; Aschwanden et al. 2001)

• Where does the mechanical energy come from?

• How is this energy coupled to the coronal plasma?

• How is the energy dissipated and converted to heat?

wavesshockseddies

(“AC”)

vs.

twistingbraiding

shear

(“DC”)vs.

reconnectionturbulenceinteract with

inhomog./nonlin.

collisions (visc, cond, resist, friction) or collisionless



Inter-granular bright points



An Alfvén wave heating model

• Cranmer & van Ballegooijen (2005) built a model of the global properties of incompressible non-WKB Alfvenic turbulence along an open flux tube.

• Background plasma properties (density, flow speed, B-field strength) were fixed empirically; wave properties were modeled with virtually no “free” parameters.

• Lower boundary condition: observed horizontal motions of G-band bright points.



MHD turbulence• It is highly likely that somewhere in the outer solar

atmosphere the fluctuations become turbulent and cascade from large to small scales:



MHD turbulence• It is highly likely that somewhere in the outer solar

atmosphere the fluctuations become turbulent and cascade from large to small scales:

• With a strong background field, it is easier to mix field lines (perp. to B) than it is to bend them (parallel to B).

• Also, the energy transport along the field is far from isotropic:

Z+Z–

Z–

(e.g., Dmitruk et al. 2002)



Turbulent heating rate

• Solid curve: predicted Qheat for a polar coronal hole.

• Dashed RGB regions: empirical estimates of heating rate of primary plasma (models tuned to match conditions at 1 AU).

• What is really needed are direct measurements of the plasma (atoms, ions, electrons) in the acceleration region of the solar wind!



UVCS / SOHO

• The Ultraviolet Coronagraph Spectrometer (UVCS) measures plasma properties of coronal protons, ions, and electrons between 1.5 and 10 solar radii.

slit field of view:• Mirror motions select height

• Instrument rolls indep. of spacecraft

• 2 UV channels: LYA & OVI

• 1 white-light polarimetry channel

• SOHO (the Solar and Heliospheric Observatory) was launched in Dec. 1995 with 12 instruments probing solar interior to outer heliosphere.

• Combines occultation with spectroscopy to do what neither alone could accomplish.



UVCS results: solar minimum (1996-1997 )

On-disk profiles: T = 1–3 million K Off-limb profiles: T > 200 million K !

• The fastest solar wind flow is expected to come from dim “coronal holes.”

• In June 1996, the first measurements of heavy ion (e.g., O+5) line emission in the extended corona revealed surprisingly wide line profiles . . .



The impact of UVCS

• The fast solar wind becomes supersonic much closer to the Sun (~2 Rs) than previously believed.

• In coronal holes, heavy ions (e.g., O+5) both flow faster and are heated hundreds of times more strongly than protons and electrons, and have anisotropic temperatures. (Kohl et al. 1997, 1998)

• At very large heights in bright streamers, the heavy ions begin to depart from thermal equilibrium in a similar way to coronal holes.

UVCS/SOHO has led to new views of the acceleration regions of the solar wind.Key results include:



Ion cyclotron waves in the corona?

• UVCS observations have rekindled theoretical efforts to understand heating and acceleration of the plasma in the (collisionless?) acceleration region of the wind.

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 phenomena

something else?



Switch gears . . .



Hot star winds: radiative driving

• Bound electron resonances have higher cross-sections than free electrons (higher “Q”)



Hot star winds: radiative driving

• Bound electron resonances have higher cross-sections than free electrons (higher “Q”)

• More acceleration facilitates more forcing!



CAK wind theory

• Highly nonlinear momentum equation has a steady-state solution only for a specific maximal mass-loss rate (“eigenvalue”); Castor, Abbott, & Klein (1975)

• Results agree well with observations.



More about pulsations . . .

• Interior: discrete spectrum of “standing waves”

• Exterior: continuous spectrum of “traveling waves”

• Nonradial pulsations (NRP) are observable at the photosphere via Doppler-shifted line profile variations:

• For hot stars, these are mainly “g-modes” (excited by deep-core convection & various opacity/ionization instabilities)



Acoustic-gravity waves: evanescence?• In an isothermal, hydrostatic atmosphere,

acoustic waves conserve energy density by growing in amplitude: δE ~ ρ (δv2)

• There is an acoustic cutoff frequency, below which waves are evanescent (non-propagating)

• Most (low-order) oscillations are evanescent in a stellar photosphere. height



Acoustic-gravity waves: evanescence?• In an isothermal, hydrostatic atmosphere,

acoustic waves conserve energy density by growing in amplitude: δE ~ ρ (δv2)

• There is an acoustic cutoff frequency, below which waves are evanescent (non-propagating)

• Most (low-order) oscillations are evanescent in a stellar photosphere.

• When a subsonic wind is considered, ALL frequencies are able to propagate!

height

v/cs = 0.01

v/cs = 0.1



Wave leakage: 1D simulation results



Stronger pulsational amplitude

• Instead of varying base density by 1% (0.99 to 1.01), vary it by a factor of 60

(i.e., 1/60 to 60), for ω/ωac ~ 0.3 :

• In the supersonic wind, acoustic waves are modified by the radiative force into the so-called “Abbott waves:” CAK critical point = sub super-Abbott flow!



Synthesized P Cygni profiles

• Profiles computed using “SEI” (Sobolev w/ Exact Integration):



Synthesized P Cygni profiles

• Profiles computed using “SEI” (Sobolev w/ Exact Integration):

• BW Vulpeculae (large-amplitude β Cep pulsator), seen with IUE:



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 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 nonadiabatic NRPs could transfer angular momentum outwards: similar to wave “radiation pressure!”



Cool stars: younger/older Sun

• How do outflows & inflows co-exist around young T Tauri stars, and does the disk accretion power the wind?

• If not, then is the wind similar to the “mature” solar wind?

• How can there be fast/supersonic winds in the chromospheres of both young & evolved solar-mass stars? waves/shocks?



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

winds of red giants and supergiants 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)



Conclusions

• Stellar winds & circumstellar waves affect a broad swath of astrophysical problems.

• Observations: spectroscopy is key!

For more information: http://cfa-www.harvard.edu/~scranmer/

• The surprisingly extreme plasma conditions in

solar coronal holes (T ion >> Tp > Te ) have guided theorists to discard some candidate processes, further investigate others, and have cross-fertilized other areas of plasma physics & astrophysics.

• Future observational programs are needed: next-generation UVCS, high-res UV stellar spectroscopy, Stellar Imager (interferometry).

More plasma diagnostics

Better understanding!



Extra slides . . .



The need for bothsolar-disk & coronagraph observations

• On-disk measurements help reveal basal coronal heating & lower boundary conditions for solar wind.

• Off-limb measurements (in solar wind “acceleration region” ) allow dynamic non-equilibrium plasma states to be followed as the asymptotic conditions at 1 AU are gradually established.

Occultation is required because extended corona is 5 to 10 orders of magnitude less bright than the disk!

Spectroscopy provides detailed plasma diagnostics that imaging

alone cannot.



Spectroscopic diagnostics• Off-limb photons formed by both collisional excitation/de-excitation and resonant

scattering of solar-disk photons.

• Profile width depends on line-of-sight component of velocity distribution (i.e., perp. temperature and projected component of wind flow speed).

• If atoms are flow in the same direction as incoming disk photons, “Doppler dimming/pumping” occurs.

• Total intensity depends on the radial component of velocity distribution (parallel temperature and main component of wind flow speed), as well as density.



Doppler dimming & pumping• After H I Lyman alpha, the O VI 1032, 1037 doublet are the next brightest lines in

the extended corona.

• The isolated 1032 line Doppler dims like Lyman alpha.

• The 1037 line is “Doppler pumped” by neighboring C II line photons when O5+ outflow speed passes 175 and 370 km/s.

• The ratio R of 1032 to 1037 intensity depends on both the bulk outflow speed (of O5+ ions) and their parallel temperature. . .

• The line widths constrain perpendicular temperature to be > 100 million K.

• R < 1 implies anisotropy!



Coronal holes: over the solar cycle• Even though large coronal holes have similar outflow speeds at 1 AU (>600

km/s), their acceleration (in O+5) in the corona is different! (Miralles et al. 2001)

Solar minimum:

Solar maximum:



Thin tubes merge into supergranular funnels

Peter (2001)

Tu et al. (2005)



Resulting wave amplitude (with damping)• Transport equations solved for 300 “monochromatic” periods (3 sec to 3 days),

then renormalized using photospheric power spectrum.

• One free parameter: base “jump amplitude” (0 to 5 km/s allowed; 3 km/s is best)



Anisotropic MHD cascade• Can MHD turbulence generate ion cyclotron waves? Many models say no!

• Simulations & analytic models predict cascade from small to large k ,leaving k ~unchanged. “Kinetic Alfven waves” with large k do not necessarily have high frequencies.

• In a low-beta plasma, KAWs are Landau-damped, heating electrons preferentially!

• Cranmer & van Ballegooijen (2003) modeled the anisotropic cascade with advection & diffusion in k-space and found some k “leakage” . . .



Streamers: open and/or closed?

• High-speed wind: strong connections to the largest coronal holes

• Low-speed wind: still no agreement on the full range of coronal sources:

hole/streamer boundary (streamer “edge”)streamer plasma sheet (“cusp/stalk”)small coronal holesactive regions (some with streamer cusps)

Wang et al. (2000)



Streamers with UVCS• Streamers viewed “edge-

on” look different in H0 and O+5

• Ion abundance depletion in “core” due to grav. settling?

• Brightest “legs” show negligible outflow, but abundances consistent with in situ slow wind.

• Higher latitudes and upper “stalk” show definite flows (Strachan et al. 2002).

• Stalk also has preferential ion heating & anisotropy, like coronal holes! (Frazin et al. 2003)



Heating above & below the critical point• Why does the critical point matter? Leer & Holzer (1980), Pneuman (1980):

• Even if heating is the same (hole vs. streamer), moving rcrit changes the above!

• Also, changing f(r) changes where the Alfven wave flux is the strongest:

• But how is an increased “Alfven wave flux” linked to actual heating?

vs.

SUBSONIC coronal heating:“puffs up” scale height,draws more particles into wind:

M u

SUPERSONIC coronal heating:subsonic region is unaffected. Energy flux has nowhere else to go:

M same, u

FA <v 2>VA Br

FA (crit) Br (crit) 1

FA Br f --------- -------- ----

Hypothesis: all flux tubes have same FA

(Kovalenko 1978; Wang & Sheeley 1991)



Why is the fast/slow wind fast/slow?

• Several ideas exist; one powerful one relates flux tube expansion to wind speed (Wang & Sheeley 1990). Physically, the geometry determines location of Parker critical point, which determines how the “available” heating affects the plasma:

vs.

SUBSONIC coronal heating:“puffs up” scale height, draws more particles into wind:

M u

SUPERSONIC coronal heating:subsonic region is unaffected. Energy flux has nowhere else to go:

M same, u

• MHD turbulence heating rates give temperatures consistent with UVCS & in situ.



CMEs• Coronal mass ejections (CMEs) are an efficient way for the Sun to shed twisted

magnetic fields (and net helicity?) and generated solar energetic particles (SEPs).

Coronagraph images contain much information, but spectroscopy also provides:

• Heating rates and energy budget (DEM)

• 3D velocity field & chirality; twisting/unwtisting rates

• conditions in shocks (preferential ion accel.)

• conditions in current sheets (reconnection rates)



UVCS CME results: Doppler shifts

Feb. 12, 2000 Intensity Width Shift(Lyman alpha) April 18, 2000



UVCS CME results: Reconnection physics• On several occasions, narrow brightening

in Fe XVIII (Te ~ 6 MK) appears in the probable location of a current sheet.

• Lin et al. (2005) also saw Lyman alpha “closing down” in the sheet: one can measure reconnection rate (Vin / Vout )



The Need for Better Observations

Even though UVCS/SOHO has made significant advances,

• We still do not understand the physical processes that heat and accelerate the entire plasma (protons, electrons, heavy ions),

• There is still controversy about whether the fast solar wind occurs primarily in dense polar plumes or in low-density inter-plume plasma,

• We still do not know how and where the various components of the variable slow solar wind are produced (e.g., “blobs”).

(Our understanding of ion cyclotron resonance is based essentially on just one ion!)

UVCS has shown that answering these questions is possible, but cannot make the required observations.

Documents

Leaky Stars: Steven R. Cranmer & many others Harvard-Smithsonian Center for Astrophysics Pulsations, Waves, and Turbulence in Stellar Winds across the