1. Turbulent Origins of the Solar Wind Steven R. Cranmer Harvard-Smithsonian Center for Astrophysics

2. Turbulent Origins of the Solar Wind Steven R. Cranmer Harvard-Smithsonian Center for Astrophysics Outline: 1. A tour of magnetic connectivity & plasma properties: wind corona chromosphere photosphere 2. Energy transport: turbulent coronal heating “recipe”

3. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 Overview: the solar atmosphere Heating is everywhere!

4. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 In situ solar wind: properties Mariner 2 (1962): first direct confirmation of continuous fast & slow solar wind. 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. By ~1990, it was clear the fast wind needs something besides gas pressure to accelerate so fast! speed (km/s) T p (10 5 K) T e (10 5 K) T ion / T p O 7+ /O 6+, Mg/O 600–800 2.4 1.0 > m ion /m p low 300–500 0.4 1.3 < m ion /m p high fastslow

5. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 In situ solar wind: connectivity 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 holes active regions Wang et al. (2000)

6. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 Coronal magnetic fields Coronal B is notoriously difficult to measure... Potential field source surface (PFSS) models have been successful in reproducing observed structures and mapping between Sun & in situ. Wang & Sheeley (1990) flux-tube expansion correlation, modified by, e.g., Arge & Pizzo (2000). Wind Speed Expansion Factor

7. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 Coronal magnetic fields: solar minimum A(r) ~ B(r) –1 ~ r 2 f(r) Banaszkiewicz et al. (1998)

8. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 Why is the fast/slow wind fast/slow? Several ideas exist; one powerful one relates the spatial dependence of the heating to the location of the Parker critical point; this determines how the “available” heating affects the plasma (e.g., Leer & Holzer 1980): 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 Banaszkiewicz et al. (1998)

9. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 Wind origins in open magnetic regions Leighton (1963) UV spectroscopy shows blueshifts in supergranular network (e.g., Hassler et al. 1999)

10. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 Supergranular “funnels” Peter (2001) Tu et al. (2005) Fisk (2005)

11. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 Granules & Supergranules

12. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 Inter-granular bright points (close-up) 100–200 km It’s widely believed that the G-band bright points are strong-field (1500 G) flux tubes surrounded by much weaker-field plasma.

13. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 Waves in thin flux tubes splitting/merging torsion longitudinal flow/wave bending (kink-mode wave) Statistics of horizontal BP motions gives power spectrum of “kink-mode” waves. BPs undergo both random walks & intermittent (reconnection?) “jumps:”

14. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 Waves in thin flux tubes splitting/merging torsion longitudinal flow/wave bending (kink-mode wave) Statistics of horizontal BP motions gives power spectrum of “kink-mode” waves. BPs undergo both random walks & intermittent (reconnection?) “jumps:” In reality, it’s not just the “pure” kink mode... (Hasan et al. 2005)

15. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 Global magnetic field connectivity Cranmer & van Ballegooijen (2005) built a model of the global properties of incompressible non-WKB Alfvenic turbulence along an open flux tube. Lower boundary condition: observed horizontal motions of G-band bright points. Along the flux tube, wave/turbulence properties should be computed consistently.

16. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 How is magnetic energy dissipated along these open flux tubes? How does this energy get into the corona to heat & accelerate the solar wind?

17. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 Coronal heating: “location location location” The basal coronal heating problem is well known: Above 2 R s, additional energy deposition is required in order to... » accelerate the fast solar wind (without artificially boosting mass loss and peak T e ), » produce the proton/electron temperatures seen in situ (also the varying magnetic moment!), » produce the strong preferential heating and temperature anisotropy of heavy ions (in the wind’s acceleration region) seen with UV spectroscopy.

18. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 UVCS/SOHO: fast solar wind 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. (e.g., Kohl et al. 1997, 1998, 2006)

19. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 Heating mechanisms A surplus of proposed ideas? (Mandrini et al. 2000; Aschwanden et al. 2001)

20. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 Heating mechanisms A surplus of proposed ideas? (Mandrini et al. 2000; Aschwanden et al. 2001) Where does the mechanical energy come from? vs.

21. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 Heating mechanisms A surplus of proposed ideas? (Mandrini et al. 2000; Aschwanden et al. 2001) Where does the mechanical energy come from? How is this energy coupled to the coronal plasma? waves shocks eddies (“AC”) vs. twisting braiding shear (“DC”) vs.

22. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 Heating mechanisms A surplus of proposed ideas? (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? waves shocks eddies (“AC”) vs. twisting braiding shear (“DC”) vs. reconnectionturbulence interact with inhomog./nonlin. collisions (visc, cond, resist, friction) or collisionless

23. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 Heating mechanisms A surplus of proposed ideas? (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? waves shocks eddies (“AC”) vs. twisting braiding shear (“DC”) vs. reconnectionturbulence interact with inhomog./nonlin. collisions (visc, cond, resist, friction) or collisionless

24. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 MHD turbulence It is highly likely that somewhere in the outer solar atmosphere the fluctuations become turbulent and cascade from large to small scales:

25. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 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–Z– Z–Z– (e.g., Matthaeus et al. 1999; Dmitruk et al. 2002)

26. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 A recipe for coronal heating? “Outer scale” correlation length (L): flux tube width (Hollweg 1986), normalized to something like 100 km at the photosphere. Z + and Z – : need to solve non-WKB Alfven wave reflection equations. Ingredients: refl. coeff = |Z + | 2 /|Z – | 2

27. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 Turbulent heating models Cranmer & van Ballegooijen (2005) solved the wave equations & derived heating rates for a fixed background state. New models: (preliminary!) self-consistent solution of waves & background one- fluid plasma state along a flux tube: photosphere to heliosphere Ingredients: Alfven waves: non-WKB reflection, turbulent damping, wave-pressure acceleration Acoustic waves: shock steepening, TdS & conductive damping, full spectrum, wave-pressure acceleration Rad. losses: transition from optically thick (LTE) to optically thin ( CHIANTI + PANDORA ) Heat conduction: transition from collisional (electron & neutral H) to collisionless “streaming”

28. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 Turbulent heating models For a polar coronal hole flux-tube: Basal acoustic flux: 10 8 erg/cm 2 /s (equiv. “piston” v = 0.3 km/s) Basal Alfvenic perpendicular amplitude: 0.4 km/s Basal turbulent scale: 120 km (G-band bright point size!) T (K) reflection coefficient Transition region is too high (8 Mm instead of 2 Mm), but otherwise not bad...

29. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 Why is the fast/slow wind fast/slow? Compare multiple 1D models in solar-minimum flux tubes with Ulysses 1st polar pass (Goldstein et al. 1996):

30. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 Why is the fast/slow wind fast/slow? Compare multiple 1D models in solar-minimum flux tubes with Ulysses 1st polar pass (Goldstein et al. 1996): “Geometry is destiny?”

31. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 Progress toward a robust recipe 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!) Probably doesn’t work for loops (keep an eye on Marco Velli) Are there additional (non-photospheric) sources of waves / turbulence / heating for open-field regions? (e.g., flux cancellation events) (B. Welsch et al. 2004) Not too bad, but...

32. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 Conclusions Theoretical advances in MHD turbulence are continuing to “feed back” into global models of the solar wind. High-resolution adaptive-optics studies of photospheric flux tubes pay off as the “bottom boundary condition” to coronal heating! More plasma diagnostics Better understanding! SOHO (especially UVCS) has led to fundamentally new views of the extended acceleration regions of the solar wind. SOHO: 1995–20?? For more information: http://cfa-www.harvard.edu/~scranmer/

33. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 Extra slides...

34. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 The solar wind 1958: Gene Parker proposed that the hot corona provides enough gas pressure to counteract gravity and accelerate a “solar wind.” 1962: Mariner 2 confirmed it! Momentum conservation: To sustain a wind,  /  t = 0, and RHS must be naturally “tuned:” Cranmer (2004), Am. J. Phys.

35. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 UVCS / SOHO 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. 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.

36. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 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...

37. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 The impact of UVCS UVCS has led to new views of the collisionless nature of solar wind acceleration. Key results include: The fast solar wind becomes supersonic much closer to the Sun (~2 R s ) 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. (e.g., Kohl et al. 1997,1998)

38. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 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.

39. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 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 O 5+ outflow speed passes 175 and 370 km/s. The ratio R of 1032 to 1037 intensity depends on both the bulk outflow speed (of O 5+ ions) and their parallel temperature... The line widths constrain perpendicular temperature to be > 100 million K. R < 1 implies anisotropy!

40. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 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, 2004) Solar minimum: Solar maximum:

41. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 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 to 10,000 Hz) suggested as a natural energy source that can be tapped to preferentially heat & accelerate heavy ions. Dissipation of these waves produces diffusion in velocity space along contours of ~constant energy in the frame moving with wave phase speed: Ion cyclotron waves in the corona lower Z/A faster diffusion

42. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 But does turbulence generate cyclotron waves? Preliminary models say “probably not” in the extended corona. (At least not in a straightforward way!) In the corona, “kinetic Alfven waves” with high k heat electrons (T >> T ) when they damp linearly. Nonlinear instabilities that locally generate high-freq. waves (Markovskii 2004)? Coupling with fast-mode waves that do cascade to high-freq. (Chandran 2006)? KAW damping leads to electron beams, further (Langmuir) turbulence, and Debye- scale electron phase space holes, which heat ions perpendicularly via “collisions” (Ergun et al. 1999; Cranmer & van Ballegooijen 2003)? How then are the ions heated & accelerated? freq. horiz. wavenumber MHD turbulence cyclotron resonance- like phenomena something else?

43. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 Alfven wave amplitude (with damping) Cranmer & van Ballegooijen (2005) solved transport equations for 300 discrete 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)

44. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 Turbulent heating rate Solid curve: predicted Q heat 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!

45. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 Streamers with UVCS Streamers viewed “edge-on” look different in H 0 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)

46. Turbulent Origins of the Solar Wind Steven R. Cranmer SHINE Workshop, July 31, 2006 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.