Turbulence and Waves as Sources for the Solar Wind

1. Turbulence and Waves as Sources for the Solar Wind Steven R. CranmerHarvard-SmithsonianCenter for Astrophysics

2. Turbulence and Waves as Sources for the Solar Wind 1. Major issues of “the debate in ’08”2. Summary of recent results3. Questions remain about both sides... Steven R. CranmerHarvard-SmithsonianCenter for Astrophysics

3. The Debate in ’08 • Two broad classes of models have evolved that attempt to self-consistently answer the question: How are fast and slow wind streams accelerated? Wave/Turbulence-Driven (WTD) models Reconnection/Loop-Opening (RLO) models My own take on the debate:arXiv: 0804.3058

4. Wave / Turbulence-Driven models • No matter the relative importance of RLO events, we do know that waves and turbulent motions are present everywhere... from photosphere to heliosphere. • How much can be accomplished by only WTD processes? (Occam’s razor?)

5. Building an Alfvén wave model • In dark intergranular lanes, strong-field photospheric flux tubes are shaken by an observed spectrum of horizontal motions. • In mainly open-field regions, Alfvén waves propagate up along the field, and partly reflect back down (non-WKB). • Nonlinear couplings allow a (mainly perpendicular) turbulent cascade, terminated by damping → gradual heating over several solar radii.

6. 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., Matthaeus et al. 1999; Dmitruk et al. 2002)

7. Self-consistent 1D models • Cranmer, van Ballegooijen, & Edgar (2007) computed solutions for the waves & background one-fluid plasma state along various flux tubes... going from the photosphere to the heliosphere. • The only free parameters: radial magnetic field & photospheric wave properties. • Ingredients: • Alfvén waves: non-WKB reflection with full spectrum, turbulent damping, wave-pressure acceleration • Acoustic waves: shock steepening, TdS & conductive damping, full spectrum, wave-pressure acceleration • Radiative losses: transition from optically thick (LTE) to optically thin (CHIANTI + PANDORA) • Heat conduction: transition from collisional (electron & neutral H) to collisionless “streaming”

8. Results: turbulent heating & acceleration T (K) Ulysses SWOOPS Goldstein et al. (1996) reflection coefficient

9. Results: frozen-in charge states Cranmer et al. (2007) WTD: Fisk (2003), Gloeckler et al. (2003) RLO: Ulysses SWICS (see also X. Wang et al. 2008) Both models need something else... coronal “halo” electrons?(Esser & Edgar 2001)

10. Results: first ionization potential effect • Cranmer et al. (2007) also showed that steady-state models can produce preferential enhancements in low-FIP elemental abundances. • Laming’s (2004) theory of Alfvén wave “ponderomotive forces” (that are dependent on ion mass, charge, and ionization potential) was assumed to be the culprit for the FIP effect. • Even though the input B-field differences were imposed high up in the extended corona, the output time-steady subsonic atmosphere (down into the upper chromosphere!) must adjust to these conditions. Ulysses SWICS

11. UVCS/SOHO Progress towards a robust WTD “recipe” Not too bad, but . . . • Because of the need to determine non-WKB (nonlocal!) reflection coefficients, it may not be easy to insert into global/3D MHD models. • Doesn’t specify proton vs. electron heating (they conduct differently) • Can MHD turbulence generate enough ion-cyclotron waves to heat heavy ions? Kohl et al. 1997, 2006; Cranmer et al. 2008; Isenberg & Vasquez (SP31D-07)

12. Peres et al. 2004 Reconnection / Loop-Opening models Some basic issues of overall “energy budget” still need to be resolved: • Do reconnections between open & closed regions cover enough of the solar surface to account for the majority of the solar wind volume? • Is the Feldman et al. (1999) scaling between loop-size and coronal temperature robust?

13. What next? • Both WTD and RLO paradigms have passed some basic “tests” of comparison with observations. What could this imply? • A combination of both ideas could work best? • Existing models don’t contain the right physics – once that is included, one or the other idea may fail to work? • Comparisons with observations haven’t been comprehensive enough to allow their true differences to be seen? (e.g., Schwadron, McGregor, Hughes) (keep plugging away at modeling...) (we need T from loop-tops to critical points!)

14. . . . not possible without standing on the shoulders of giants! Conclusions • Theoretical advances in MHD turbulence continue to “feed back” into global models of the solar wind, as well as into many other areas of plasma physics and astrophysics. • The debate between waves/turbulence and reconnection/loop-opening mechanisms of solar wind acceleration goes on . . . vs. For more information:arXiv: 0804.3058

15. Extra slides . . .

16. The Debate in ’08 • Two broad classes of models have evolved that attempt to self-consistently answer the question: How are fast and slow wind streams accelerated? Wave/Turbulence-Driven (WTD) models Reconnection/Loop-Opening (RLO) models My own take on the debate:arXiv: 0804.3058

17. The extended solar atmosphere . . . Heating is everywhere . . . . . . and everything is in motion

18. fast slow speed (km/s) Tp(105 K) Te(105 K) Tion / Tp O7+/O6+, Mg/O 600–800 2.4 1.0 > mion/mp low 300–500 0.4 1.3 < mion/mp high In situ solar wind: properties • Mariner 2 detected two phases of solar wind: slow (mostly) + fast streams • Uncertainties about which type is “ambient” persisted because measurements were limited to the ecliptic plane . . . • Ulysses left the ecliptic; provided 3D view of the wind’s source regions. • Helios saw strong departures from Maxwellians. By ~1990, it was clear the fast wind needs something besides gas pressure to accelerate so fast!

19. Intensity modulations . . . • Motion tracking in images . . . • Doppler shifts . . . • Doppler broadening . . . • Radio sounding . . . Waves: remote-sensing techniques The following techniques are direct… (UVCS ion heating was more indirect) Tomczyk et al. (2007)

20. Wang et al. (2000) Solar wind: connectivity to the corona • High-speed wind: strong connections to the largest coronal holes hole/streamer boundary (streamer edge) streamer plasma sheet (“cusp/stalk”) small coronal holes active regions • Low-speed wind: still no agreement on the full range of coronal sources:

21. The coronal heating problem • We still don’t understand the physical processes responsible for heating up the coronal plasma. A lot 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 in tiny twisted & braided magnetic flux tubes. 3. Collisions (particle-particle? wave-particle?) release energy as heat. Heating Solar wind acceleration!

22. Coronal heating mechanisms • So many ideas, taxonomy is needed! (Mandrini et al. 2000; Aschwanden et al. 2001) • Where does the mechanical energy come from? • How rapidly 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

23. Reconnection / Loop-Opening models • There is a natural appeal to the RLO idea, since only a small fraction of the Sun’s magnetic flux is open. Open flux tubes are always near closed loops! • The “magnetic carpet” is continuously churning . . . • Open-field regions show coronal jets (powered by reconnection?) that contribute to the wind mass flux. Fisk (2005) Hinode/XRT (X-ray) http://xrt.cfa.harvard.edu STEREO/EUVI (195 Å) courtesy S. Patsourakos

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

25. Particles are not in “thermal equilibrium” …especially in the high-speed wind. mag. field WIND at 1 AU (Steinberg et al. 1996) Helios at 0.3 AU (e.g., Marsch et al. 1982) WIND at 1 AU (Collier et al. 1996)

26. Mirror motions select height • UVCS “rolls” independently of spacecraft • 2 UV channels: • 1 white-light polarimetry channel LYA (120–135 nm) OVI (95–120 nm + 2nd ord.) The UVCS instrument on SOHO • 1979–1995: Rocket flights and Shuttle-deployed Spartan 201 laid groundwork. • 1996–present: 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 reveal the solar wind acceleration region! slit field of view:

27. Emission lines as plasma diagnostics • Many of the lines seen by UVCS are formed by resonantly scattered disk photons. • If profiles are Doppler shifted up or down in wavelength (from the known rest wavelength), this indicates the bulk flow speed along the line-of-sight. • The widths of the profiles tell us about random motions along the line-of-sight (i.e., temperature) • The total intensity (i.e., number of photons) tells us mainly about the density of atoms, but for resonant scattering there’s also another “hidden” Doppler effect that tells us about the flow speedsperpendicular to the line-of-sight. • If atoms are flow in the same direction as incoming disk photons, “Doppler dimming/pumping” occurs.

28. 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!

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

31. 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 somek “leakage” . . .

32. So does turbulence generate cyclotron waves? Directly from the linear waves? Probably not! How then are the ions heated and accelerated? • When MHD turbulence cascades to small perpendicular scales, the small-scale shearing motions may be able to generate ion cyclotron waves (Markovskii et al. 2006). • If MHD turbulence exists for both Alfvén and fast-mode waves, the two types of waves can nonlinearly couple with one another to produce high-frequency ion cyclotron waves (Chandran 2006). • If nanoflare-like reconnection events in the low corona are frequent enough, they may fill the extended corona with electron beams that would become unstable and produce ion cyclotron waves (Markovskii 2007). • If kinetic Alfvén waves reach large enough amplitudes, they can damp via wave-particle interactions and heat ions (Voitenko & Goossens 2006; Wu & Yang 2007). • Kinetic Alfvén wave damping in the extended corona could lead to electron beams, Langmuir turbulence, and Debye-scale electron phase space holes which heat ions perpendicularly via “collisions” (Ergun et al. 1999; Cranmer & van Ballegooijen 2003).

33. Future diagnostics: more spectral lines! • How/where do plasma fluctuations drive the preferential ion heating and acceleration, and how are the fluctuations produced and damped? • Observing emission lines of additional ions (i.e., more charge & mass combinations) would constrain the specific kinds of waves and the specific collisionless damping modes. Comparison of predictions of UV line widths for ion cyclotron heating in 2 extreme limits (which UVCS observations [black circles] cannot distinguish). Cranmer (2002), astro-ph/0209301

34. Future Diagnostics: electron VDF • Simulated H I Lyman alpha broadening from both H0 motions (yellow) and electron Thomson scattering (green). Both proton and electron temperatures can be measured.

35. Synergy with other systems • T Tauri stars: observations suggest a “polar wind” that scales with the mass accretion rate. Cranmer et al. (2007) code is being adapted to these systems... • Pulsating variables: Pulsations “leak” outwards as non-WKB waves and shock-trains. New insights from solar wave-reflection theory are being extended. • AGN accretion flows: A similarly collisionless (but pressure-dominated) plasma undergoing anisotropic MHD cascade, kinetic wave-particle interactions, etc. Freytag et al. (2002) Matt & Pudritz (2005)