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Solar/Stellar Winds: An Overview

Solar/Stellar Winds: An Overview. (Untapped gold mines for space physicists?). Steven R. Cranmer Harvard-Smithsonian Center for Astrophysics. Outline. The Stellar Zoo:. The Sun Young low-mass stars (T Tauri, FU Ori, etc.) Evolved low-mass stars (giants, supergiants, Miras)

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Solar/Stellar Winds: An Overview

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

  2. 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?

  3. 106 104 102 1 10–2 no coronae? I "cool" dense (slow?) winds radiatively driven winds "warm" hybrid winds III Be stars "hot" solar-type winds V Sun  flare stars M O B A F G K 30,000 10,000 6,000 3,000 Stellar winds on the HR Diagram

  4. ion opacity? (Teff > 15,000 K) • free electron (Thomson) opacity? (goes as 1/r2 ; needs to be supplemented) • dust opacity? (Teff < 3,500 K) ~ ~ 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 . . . F = ma

  5. 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 . . .

  6. 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

  7. The solar wind acceleration debate • What determines how much energy and momentum goes into the solar wind? • Waves & turbulence input from below? • Reconnection & mass input from loops? vs.

  8. The solar wind acceleration debate • What determines how much energy and momentum goes into the solar wind? • Waves & turbulence input from below? • Reconnection & mass input from loops? vs. • 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.

  9. 106 104 102 1 10–2 no coronae? I "cool" dense (slow?) winds radiatively driven winds "warm" hybrid winds III Be stars "hot" solar-type winds V Sun  flare stars M O B A F G K 30,000 10,000 6,000 3,000 Stellar winds on the HR Diagram

  10. Pre-Main-Sequence accretion phases Feigelson & Montmerle (1999) M. Burton (UNSW)

  11. T Tauri stars: accreting young Suns • 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. (Romanova et al. 2007) (Matt & Pudritz 2005, 2008)

  12. 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!

  13. Stellar winds on the HR Diagram 106 104 102 1 10–2 no coronae? I "cool" dense (slow?) winds radiatively driven winds "warm" hybrid winds III Be stars "hot" solar-type winds V Sun flare stars M O B A F G K 30,000 10,000 6,000 3,000

  14. Cool-star mass loss rates

  15. Cool-star dimensional analysis . . . • Stellar wind power: • Reimers (1975, 1977) proposed a semi-empirical scaling: • 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.)

  16. Cool-star mass loss rates Schröder & Cuntz (2005) scaling for lum. classes I, III, V

  17. 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.

  18. 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)

  19. 106 104 102 1 10–2 no coronae? I "cool" dense (slow?) winds radiatively driven winds "warm" hybrid winds III Be stars "hot" solar-type winds V Sun  flare stars M O B A F G K 30,000 10,000 6,000 3,000 Stellar winds on the HR Diagram

  20. 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. • 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. Kudritzki et al.

  21. 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. ? MPG courtesy John Telting Massa 1994

  22. 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. • Collisions of parcels lead to stochastic upflows/downflows. • Fast rotation: a rigid “magnetosphere” develops (seen in Ap/Bp stars?) faster rotation stronger field

  23. Conclusions A goal of this round of NESSC: • Positive feedback and collaboration between the solar physics, space physics, plasma physics, and astrophysics communities. • What do you find interesting? Topics for discussion (later?): • “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

  24. Extra slides . . .

  25. Diagnostics of stellar mass loss • 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 • Radio: free-free emission from (partially ionized?) components of the wind • Continuum methods need V from another diagnostic to get mass loss rate. • Clumping? (Bernat 1976) wind star (van den Oord & Doyle 1997)

  26. 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)

  27. 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)

  28. 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 ~ LX1.3 M ~ LX1.0 M ~ LX0.4 M ~ LX0.1

  29. Multi-fluid collisionless effects? Polar coronal hole model

  30. Multi-fluid collisionless effects? O+5 O+6 protons electrons (thermal core only)

  31. Alfven wave’s oscillating E and B fields ion’s Larmor motion around radial B-field something else? Departures from thermal equilibrium • UVCS/SOHO observations rekindled theoretical efforts to understand collisionless heating and acceleration effects in the extended corona. • Ion cyclotron waves (10–10,000 Hz) suggested as a “natural” energy source that can be tapped to preferentially heat & accelerate heavy ions. cyclotron resonance-like phenomena MHD turbulence

  32. 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) • 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. Rot. Period (days) L. Hartmann, lecture notes

  33. 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)

  34. 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.)

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