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Leaky Stars:

Leaky Stars:. Pulsations, Waves, and Turbulence in Stellar Winds across the H-R Diagram. Steven R. Cranmer & many others Harvard-Smithsonian Center for Astrophysics. Leaky Stars:. Outline: Background: history & basic physics The Sun: coronal heating & fast solar wind

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Leaky Stars:

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  1. Leaky Stars: Pulsations, Waves, and Turbulencein Stellar Windsacross the H-R Diagram Steven R. Cranmer & many othersHarvard-Smithsonian Center for Astrophysics

  2. Leaky Stars: • 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? Pulsations, Waves, and Turbulencein Stellar Windsacross the H-R Diagram Steven R. Cranmer & many othersHarvard-Smithsonian Center for Astrophysics

  3. Motivations . . . Solar corona & wind: • “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. Stellar winds: • 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?

  4. “New stars” 1572: Tycho’s supernova 1600: P Cygni outburst (“Revenante of the Swan”) 1604: Kepler’s supernova in “Serepentarius” First observations of stellar outflows ? • Coronae & Aurorae seen since antiquity . . .

  5. 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)

  6. Brief history: stellar winds • 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) O supergiant (Morton 1967) M supergiant (Bernat 1976) • Also: IR excesses, maser emission, “plain” blueshifts.

  7. Schematic H-R Diagram 106 104 102 1 10–2 I III V Sun M O B A F G K 30,000 10,000 6,000 3,000

  8. Stellar winds 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

  9. Convection zones 106 104 102 1 10–2 I deep core convection fully convective III V Sun subsurface convection M O B A F G K 30,000 10,000 6,000 3,000

  10. Temperatures in outer atmospheres Sun Hot star (O, B) Cool star (K, M) ?

  11. One-page stellar wind physics • Momentum conservation: To sustain a wind, /t = 0 , and RHS must be “tuned:”

  12. One-page stellar wind physics • Momentum conservation: To sustain a wind, /t = 0 , and RHS must be “tuned:” • Energy conservation:

  13. 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)

  14. One-page stellar wind physics • Momentum conservation: To sustain a wind, /t = 0 , and RHS must be “tuned:” • Energy conservation: • Chromosphere: heating  rad. losses • Photosphere (& most of hot-star wind)

  15. One-page stellar wind physics • Momentum conservation: To sustain a wind, /t = 0 , and RHS must be “tuned:” • Energy conservation: • Transition region & low corona • Chromosphere: heating  rad. losses • Photosphere (& most of hot-star wind)

  16. One-page stellar wind physics • Momentum conservation: To sustain a wind, /t = 0 , and RHS must be “tuned:” • Energy conservation: • Extended corona & cool-star wind • Transition region & low corona • Chromosphere: heating  rad. losses • Photosphere (& most of hot-star wind)

  17. 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 . . .

  18. 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 . . .

  19. 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? vs. waves shocks eddies (“AC”) twisting braiding shear (“DC”) vs. interact with inhomog./nonlin. turbulence reconnection collisions (visc, cond, resist, friction) or collisionless

  20. 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? vs. waves shocks eddies (“AC”) twisting braiding shear (“DC”) vs. interact with inhomog./nonlin. turbulence reconnection collisions (visc, cond, resist, friction) or collisionless

  21. Inter-granular bright points

  22. 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.

  23. MHD turbulence • It is highly likely that somewhere in the outer solar atmosphere the fluctuations become turbulent and cascade from large to small scales:

  24. 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)

  25. 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!

  26. UVCS / SOHO • SOHO (the Solar and Heliospheric Observatory) was launched in Dec. 1995 with 12 instruments probing solar interior to outer heliosphere. • The Ultraviolet Coronagraph Spectrometer (UVCS) measures plasma properties of coronal protons, ions, and electrons between 1.5 and 10 solar radii. • Combines occultation with spectroscopy to do what neither alone could accomplish. slit field of view: • Mirror motions select height • Instrument rolls indep. of spacecraft • 2 UV channels: LYA & OVI • 1 white-light polarimetry channel

  27. On-disk profiles: T = 1–3 million K Off-limb profiles: T > 200 million K ! UVCS results: solar minimum (1996-1997) • 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 . . .

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

  29. Alfven wave’s oscillating E and B fields ion’s Larmor motion around radial B-field something else? 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. • 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

  30. Switch gears . . .

  31. Hot star winds: radiative driving • Bound electron resonances have higher cross-sections than free electrons (higher “Q”)

  32. Hot star winds: radiative driving • Bound electron resonances have higher cross-sections than free electrons (higher “Q”) • More acceleration facilitates more forcing!

  33. 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.

  34. 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)

  35. 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

  36. 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 v/cs = 0.01 v/cs = 0.1 • When a subsonic wind is considered, ALL frequencies are able to propagate!

  37. Wave leakage: 1D simulation results

  38. 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!

  39. Synthesized P Cygni profiles • Profiles computed using “SEI” (Sobolev w/ Exact Integration):

  40. Synthesized P Cygni profiles • Profiles computed using “SEI” (Sobolev w/ Exact Integration): • BW Vulpeculae (large-amplitude β Cep pulsator), seen with IUE:

  41. 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!”

  42. 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?

  43. 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)

  44. More plasma diagnostics Better understanding! Conclusions • Stellar winds & circumstellar waves affect a broad swath of astrophysical problems. • Observations: spectroscopy is key! • 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). For more information: http://cfa-www.harvard.edu/~scranmer/

  45. Extra slides . . .

  46. 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.

  47. 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). • Total intensity depends on the radial component of velocity distribution (parallel temperature and main component of wind flow speed), as well as density. • If atoms are flow in the same direction as incoming disk photons, “Doppler dimming/pumping” occurs.

  48. 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!

  49. 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:

  50. Thin tubes merge into supergranular funnels Peter (2001) Tu et al. (2005)

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