2. U.W., April 14, 2005 Solar flares in the new millennium H.S. Hudson SSL/UCB

3. U.W., April 14, 2005 Outline Physical and historical background The nature of the corona Current problems in flare research Flares, especially as viewed by RHESSI Conclusions Miscellaneous RHESSI things

4. U.W., April 14, 2005 Conclusions New things this millennium: CMEs, helicity,  -rays RHESSI is allowing us to understand the dominant role of accelerated particles Theoretically, we are forced to go beyond MHD theory and magnetic reconnection

5. U.W., April 14, 2005

6. Two missing links: Heaviside and Rontgen

7. U.W., April 14, 2005

8. “Sudden Flare Effect” Compass deflections result from enhanced ionospheric current system Recent suggestion that the ionization causing this results from  -rays, not soft X- rays Echos of the physics involved in magnetar and lightning  -ray behavior

9. U.W., April 14, 2005 TRACE EUV observations Issues…

10. U.W., April 14, 2005 TRACE 1600A TRACE 195A Shrinkage, dimming, oscillation

11. U.W., April 14, 2005 G. A. Gary, Solar Phys. 203, 71 (2001) (v A ~ 200  -1/2 km/s) CH Distribution of coronal plasma 

12. U.W., April 14, 2005 Field and energy are concentrated in active regions Active-region magnetic fields via Roumeliotis-Wheatland technique (McTiernan) Mass loading via empirical law (Lundquist/Fisher)

13. U.W., April 14, 2005 NOAA 10486, Haleakala IVM data,  cube Roumeliotis-Wheatland-McTiernan method 64x64x64x ~3000 km ScaledNot scaled

14. U.W., April 14, 2005 Lundquist et al., SPD 2004

15. U.W., April 14, 2005 Summing up the corona It’s like a spherical condenser filled with a low-  dielectric (about 700  F?) The upper boundary is the solar wind, which is massive The lower boundary (the “transition layer”) is extraordinarily complex and not at all understood yet Mysterious things happen: flares and CMEs

16. U.W., April 14, 2005 Coronal Mass Ejection

17. U.W., April 14, 2005 What is the mass of the solar wind? How big is a CME? M corona = 1.3 x 10 18 g (1-10 R sun ) M sw = 0.7 x 10 18 g (10-100 R sun ) M heliosphere = 7 x 10 18 g (100-1000 R sun ) M infinity = infinite g (universe)  CME = 0.13 sr (40 o FWHM) Withbroe (1988) “quiet corona” model with 1/r 2 extension

18. U.W., April 14, 2005 Current problem areas How does magnetic energy penetrate the corona/photosphere boundary? Why does coronal magnetic reconnection not readily happen? How does particle acceleration work in solar flares?

19. U.W., April 14, 2005 Emslie et al., JGR (2004) Hudson 2005?

20. U.W., April 14, 2005 Major breakthrough! (Woods et al., GRL 2004) First bolometric observations of a solar flare (SORCE satellite) Detection of the impulsive phase Background noise essentially from the p-modes ~ 300  mag

21. U.W., April 14, 2005 RHESSI Particle acceleration is key to understanding RHESSI can image not only hard X-ray sources, but  -rays as well RHESSI has extraordinary sensitivity

22. U.W., April 14, 2005 Reuven Ramaty 1937 – 2001

23. U.W., April 14, 2005 SPECTROSCOPY π 0 Decay Nonthermal Bremsstrahlung Thermal Bremsstrahlung Composite Solar Flare Spectrum Positron and Nuclear Gamma-Ray lines T = 2 x 10 7 K T = 4 x 10 7 K Fe

24. U.W., April 14, 2005 10-100 keV electrons The Neupert effect The soft-hard-soft spectral pattern “Escape” into the heliosphere

25. U.W., April 14, 2005 The Neupert effect

26. U.W., April 14, 2005 The soft-hard-soft pattern

27. U.W., April 14, 2005 Flare image morphology Ribbons and footpoints in hard X-rays Conjugacy Gamma-ray sources

28. U.W., April 14, 2005 EUV flare ribbons and hard X-ray footpoint sources

29. U.W., April 14, 2005 Gamma-ray imaging too…

30. U.W., April 14, 2005 Problems Why are the hard X-ray footpoints so compact, when the ribbons are extended? Why are the  -ray sources displaced from the X-ray sources? …prosaic reasons? Interesting reasons?

31. U.W., April 14, 2005 A major problem: bremsstrahlung is inefficient, and in a major event we need as many as 10 36 e - /s. But the footpoint areas are small!

32. U.W., April 14, 2005 IR 1.56m observations (Xu et al., 2004; should show the opacity minimum height)

33. U.W., April 14, 2005 Hard X-ray footpoint behavior (S. Krucker)

34. U.W., April 14, 2005 Velocity vs. rate of energy loss of electrons Nine intervals, nine spectra; thick target model  energy deposition (rate of energy loss by non- thermal electrons in footpoints) 2 Reconnection rate d F /dt= B v a v= velocity B= magnetic field strength a=footpoint diameter B hard to observe for near limb flare B~1000 G; a~2000km  d F /dt ~ 2e18 Mx/s  E ~ 5 kV/m Higher rate of reconnection produces more energetic electrons per unit time

35. U.W., April 14, 2005 E time RHESSI Soft X-rays

36. U.W., April 14, 2005 Status of flare theory Standard model of large-scale magnetic reconnection… Cartoons illustrating this (and other) models http://solarmuri.ssl.berkeley.edu/~hhudson/cartoons/ Have major uncertainties for - Coupling of scales - Particle acceleration - Role of helicity

37. U.W., April 14, 2005

38. ‘Somewhat analogous outbursts often happen on the Sun, in explosive events called "solar flares." During a solar flare, magnetic field lines near the Sun's surface change the pattern by which they connect to each other, a process called "magnetic reconnection" which releases pure magnetic energy. This happens in magnetars too.’ http://solomon.as.utexas.edu/~duncan/magnetar.html Introduction to flare theory via the physics of magnetars

39. U.W., April 14, 2005 RHESSI magnetar response

40. U.W., April 14, 2005 Middle-aged star decade slowly solar flare

41. U.W., April 14, 2005

43. Magnetic reconnection “Reconnection” is only descriptive and does not describe the physics Magnetic restructuring is necessary for flare energy release from the magnetic field Clear unambiguous evidence is hard to find

44. U.W., April 14, 2005 Gosling et al., JGR 110, A01107, 2005

45. U.W., April 14, 2005 Conclusions New things this millennium: CMEs, helicity,  -rays RHESSI is allowing us to understand the dominant role of accelerated particles Theoretically, we are forced to go beyond MHD theory and magnetic reconnection

46. U.W., April 14, 2005 END

47. U.W., April 14, 2005 Movie of dimming (Aug 28, 1992) Coronal Dimming

48. U.W., April 14, 2005 Moreton-Ramsey wave and EIT wave Thompson et al., 1998

49. U.W., April 14, 2005 Type III (“fast drift”) Type II (“slow drift”), harmonic “Ignition”

50. U.W., April 14, 2005 Inference of source motion in drifting radio bursts Assume a density model  (z) with height z, normally empirical (e.g. “fourfold Newkirk”) Determine the drift rate in MHz/s Convert to height from plasma-frequency assumption; f p = 9000 n e 0.5 Hz Typically, assume radial motion

51. U.W., April 14, 2005 Cartoons illustrating wave origins? cf. http://solarmuri.ssl.berkeley.edu/~hhudson/cartoons There doesn’t seem to be a satisfactory cartoon! Sturrock CME Hudson flare

52. U.W., April 14, 2005 The CME-driven shock in the corona The CME involves outward plasma motions perpendicular to the field We see the result of these motions as dimmings, but the data are not good enough to follow the flows nor to see a bow wave There is an Alfven-speed “hole” in the middle corona in which Mach numbers could be larger

53. U.W., April 14, 2005 SUMMARY Coronal shock waves (metric type II) are blast waves (Uchida) launched by compact structures at flare onset. These propagate in an undisturbed corona The CME eruption restructures the corona and pushes a bow wave ahead of it into the solar wind. This creates a type II burst at long wavelengths

54. U.W., April 14, 2005 Where do “Solar cosmic rays” come from? Consensus holds that CME-driven shocks are responsible for most SEPs, but that something else is also happening Shock geometry and Mach numbers in the high corona are crucial factors: quasi-perpendicular fronts and large Mach numbers preferred The theory is incomplete but PIC simulations are appearing for the planetary bow shocks, at least

55. U.W., April 14, 2005 Imaging of coronal shocks: good news and bad news A shock wave should provide a sharp density gradient, easy to detect in images We can observe motions in two dimensions The medium is optically thin => confusion The wave may not be bright compared with other flare components The corona generally has low plasma beta, so the observed mass may not be structurally important

56. U.W., April 14, 2005 … Only imaging can properly characterize the large-scale structure The solar corona isn’t really accessible any other way

57. U.W., April 14, 2005 Imaging of coronal shocks Type II bursts (plasma radiation) Moreton waves (H  in the chromosphere) New modalities: EIT, X-rays 1, microwaves, He 10830, meter waves (thermal), meter waves (nonthermal) 1 Three events: Khan & Aurass (2002); Narukage et al. (2002); Hudson et al. (2003)

58. U.W., April 14, 2005 Direct X-ray observation Uchida 1968 Yohkoh 1998 EIT

59. U.W., April 14, 2005 Why X-ray waves are hard to observe directly Pre-flare transect Flare transect The wave - just ripples on the scattering halo!

60. U.W., April 14, 2005 Heliospheric shocks in images? Maia et al., ApJ 528, L49 (2000) Vourlidas et al., ApJ 598, 1392 (2003) SOHO/UVCS

61. U.W., April 14, 2005 Vourlidas et al., ApJ 598, 1392 (2003) Where is the bow shock ?

62. U.W., April 14, 2005 Inferring the Mach number Method: Estimate temperature jump from soft X-ray images and apply Rankine-Hugoniot condition

63. U.W., April 14, 2005 X-ray signal S ~ n e 2 f(T) f(T) ~ T 2  ln(S)/  ln(n) ~ 2  Mach number estimate for 6 May 1998 event