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Combined WIND-RHESSI-TRACE studies

This study combines observations from WIND, RHESSI, and TRACE to investigate the origin of solar impulsive energetic electrons. It examines the escape of electrons from the Sun, their acceleration in flares and shocks, and the timing and propagation effects of these events. The study also explores the use of multi-point measurements and radio tracking for coordinating observations.

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Combined WIND-RHESSI-TRACE studies

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  1. Combined WIND-RHESSI-TRACE studies . . Sam Krucker Space Sciences Lab, UC Berkeley The origin of solar impulsive energetic electrons: combined observations

  2. Impulsive electron events observed at 1 AU WIND/3DP: ~15 event/year during solar min. STEREO/STE: 50 times more sensitive below 30 keV  more events STEREO first 2 years

  3. STEREO/STE observations • Solid state detectors down to 2 keV (wind/3dp: 30 keV) • 50 times more sensitive  5 times more events • Accurate onset times down to 2 keV • 2 point measurements + WIND

  4. ? energetic electrons escaping from the Sun e- Sun EM radiation magnetic field line REMOTE sensing observations IN-SITU electron observations STEREO, WIND, ACE

  5. FLARE accelerated electrons escape flare e- Sun magnetic field line IN-SITU electron observations STEREO, WIND, ACE

  6. FLARE accelerated electrons escape flare e- Sun magnetic field line e- IN-SITU electron observations STEREO, WIND, ACE

  7. Escaping electrons could also be accelerated late in the flare  no correlation with impulsive phase flare e- Sun acc. site of escaping electrons EM radiation magnetic field line main flare REMOTE sensing observations IN-SITU electron observations STEREO, WIND, ACE HXR footpoints

  8. SHOCK accelerated electrons flare e- shock Flare accelerated electrons do not escape or escape along field lines not connected to the spacecraft magnetic field line IN-SITU electron observations STEREO, WIND, ACE

  9. Flare or shock acceleration? • Different timing • Depending of magnetic connection different component are observed

  10. Timing From onset times at 1 AU (velocity dispersion) solar release time can be approximated. Controversy: propagation effect or scattering? Slope gives path length Intersection gives release time 30 keV

  11. Timing From onset times at 1 AU (velocity dispersion) solar release time can be approximated. Controversy: propagation effect or scattering? WIND/3DP Electrostatic analyzers: large error bars Solid state detectors: SMALL error bars (~few minutes) 30 keV

  12. Timing From onset times at 1 AU (velocity dispersion) solar release time can be approximated. Controversy: propagation effect or scattering? STEREO/STE Solid state detectors down to 2 keV STEREO WIND/3DP observerations Electrostatic analyzers: large error bars Solid state detectors: SMALL error bars (~few minutes)

  13. What is reported? flare e- shock • Same timing as flare for some events magnetic field line e.g. Krucker et al. 1999, Maia & Pick 2004, Klein et al. 2005 What can be done with multi-point measurements?

  14. What is reported? e- flare shock • Same timing as flare for some events • >30 keV electron often delayed  shock magnetic field line e.g. Krucker et al. 1999, Haggerty & Roelof 2002, Maia & Pick 2004, Klein et al. 2005 What can be done with multi-point measurements?

  15. What is reported? e- flare e- shock • Same timing as flare for some events • >30 keV electron often delayed  shock • <20 keV more often with flare magnetic field line Wang et al. 2006 What can be done with multi-point measurements?

  16. What is reported? e- flare e- shock • Same timing as flare for some events • >30 keV electron often delayed  shock • <20 keV more often with flare • timing different because of scattering  all events are flare related magnetic field line Cane 2004 What can be done with multi-point measurements?

  17. earlier shock magnetic field line e- Early on, only one spacecraft is connected to the shock

  18. earlier later e- shock shock magnetic field line e- e- Early on, only one spacecraft is connected to the shock Later both spacecrafts are connected to shock  different onset times are expected

  19. Shock accelerated electrons are seen by STEREO 2 flare e- shock e- Flare accelerated electrons are seen by STEREO 1 magnetic field line STEREO 2: Later onset STEREO 1: Earlier on set expected

  20. Timing alone not conclusive.Combination with imaging and modeling needed!

  21. Escaping electrons produce type III bursts e- Acceleration site e- • Electrons lose their • energy by collisions • X-ray emission • heating heating

  22. Coronal imaging EUV/X-ray observations reveal coronal structures. STEREO: 3d structure, SOLAR B: X-rays, B, flows, RHESSI: HXRs Sun jet e-

  23. movie

  24. Coronal imaging EUV/X-ray observations reveal coronal structures. STEREO: 3d structure, SOLAR B: X-rays, B, flows, RHESSI: HXRs Sun jet e- What are chances to observed an event? How to coordinate observations?

  25. Radio tracking 400-150 MHz: NRH In the future: FASR type III bursts (electron beams) Simulated radio positions 1-2 solar radii <16 MHz: STEREO/WAVES open field line

  26. Radio tracking (K.-L. Klein): potential magnetic field extrapolation (Schrijver & Derosa 2003) N N type III 432 MHz 327 MHz 164 MHz to Earth S May 1, 2000 type III Only open field lines are plotted.

  27. Type II burst: RADIO TRACKING gives shock location flare e- shock e- Type III burst: RADIO TRACKING gives path of electron beam magnetic field line Compare with onset times & 3D observations & modeling

  28. Type II burst: RADIO TRACKING gives shock location flare e- shock e- What is possible? Type III burst: RADIO TRACKING gives path of electron beam magnetic field line Compare with onset times & 3D observations & modeling

  29. Summary • Combined observations have great potential • Timing studies combined with imaging and modeling

  30. Comparing spectra PHOTON SPECTRA: Produced by downward moving electron beam ELECTRON SPECTRA: spectrum of escaping electrons  rough correlation d g

  31. Comparing spectra PHOTON SPECTRA: Power law fit to HXR spectra averaged over peak ELECTRON SPECTRA: Power law fit to peak flux • Assuming power spectra: • THIN: d = g – 1 • THICK: d = g + 1 • RESULTS: • correlation seen • values are between d g

  32. STEREO 2 is not connected to flare site flare e- Flare accelerated electrons are seen by STEREO 1 magnetic field line STEREO 2: No particles seen STEREO 1  Better estimates of total number of electrons (energy)

  33. Electron spectrum at 1AU Typical electron spectrum can be fitted with broken power law: Break around: 30-100 keV Steeper at higher energies Oakley, Krucker, & Lin 2006

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