Group V: Report Regular Members: K. Arzner, A. Benz, C. Dauphin, G. Emslie, M. Onofri, N. Vilmer, L. Vlahos Visitors: E. Kontar, G. Mann, R. Lin, V. Zharkova

Main Goals 1.Constrains on particle acceleration from the RHESSI data (close collaboration with all WGs) and other available sources of information on high energy particles 2.Discuss new theories on particle acceleration 3.Connecting theories on particle acceleration with the global magnetic topologies hosting flares and CMEs

Constraints on Acceleration/Transport(Electrons) Must produce an electron flux of at least electrons per second Must be able to accelerate electrons on time scales at most 10 milliseconds Must sometimes produce electron energies greater than at least 10’s of GeV Mechanism must be able to produce a flattening of the electron distribution at energies on the order of 500 keV Higher nonthermal hard X-ray flux statistically associated with harder spectra

The Electron “Problem” Efficiency of bremsstrahlung production ~ (ergs of X-rays per erg of electrons)  Electron flux ~ 10 5  hard X-ray flux Electron energy can be – ergs in large events Total number of accelerated electrons up to (cf. number of electrons in loop ~10 38 ). –replenishment and current closure necessary

Revised Numbers ModeSymbolLog (Energy) April 21, 2002July 23, 2002 MagneticUBUB 32.3 ± 0.3 Flare Intermediate ThermalU th 31.3 (+0.4,-1)31.1 (+0.4,-1) ElectronsUeUe 31.3 (+?, -0.5)31.5 (+?, -0.5) IonsUi< ± 0.5 Final SXR RadiationURUR Total RadiationURUR > 31.7> 31.6 CME KineticUKUK 32.3 ± ± 0.3 PotentialUU 30.7 ± ± 0.3 SEPsUPUP 31.5 ± 0.6< 30

X/  -ray spectrum RHESSI Energy range Pion decay radiation (ions > 100 MeV/nuc) sometimes with neutrons Ultrarelativistic Electron Bremsstrahlung Thermal components Electron bremsstrahlung  -ray lines (ions > 3 MeV/nuc) T= K T= K SMM/GRS Phebus/Granat Observations GAMMA1 GRO GONG

Electron-Dominated Events First observed with SMM (Rieger et al, 1993) Short duration (s to 10 s) high energy (> 10 MeV) bremsstrahlung emission No detectable GRL flux Photon spectrum > 1 MeV (  X  -1.5—2.0) For 2 PHEBUS events o if W i>1MeV/nuc  W e>20 keV oNo detectable GRL above continuum oWeak GRL flares? Vilmer et al (1999) PHEBUS BATSE

thermal non-thermal RHESSI two component fits: T, EM γ, F 35

Grigis & B. flux spectral index

Energy dependent photon spectral index Interval 3 (peak of the flare) Spectral index evolution:

Mean Electron Spectrum: Temporal evolution Temporal evolution of the Regularized Mean Electron Spectrum (20s time intervals) RHESSI Lightcurves 3-12keV; 12-25keV; 25-50keV; keV

GOES 1-8 A DERIVATIVE Non-thermal preflare coronal sources

RHESSI SPECTRA 5-50 keV Thermal+broken powerlaw Preflare period: 01:02:00-01:11:00 Broken powerlaw extends down to 5 keV Thermal component never dominates EM and T are poorly determined Chisquare ~ 1 if EM=0 White = photons, Green = thermal model, Red = broken powerlaw, Purple = background (NB similar source in July 23 rd 2002 event)

Electron spectrum at 1AU Typical electron spectrum can be fitted with broken power law: Break around: keV Steeper at higher energies Oakley, Krucker, & Lin 2004

Ions Tens of MeV ions and hundreds of MeV particles can be accelerated at the same time; We also see cases where we see a stage when hundreds of MeV ions are primarily accelerated.

| sec----| | sec------| keV keV 4 – 6.4 MeV  -ray line emission can be delayed from hard X-rays from <2 to 10’s of sec.

June 3, Evidence for delayed high-energy emission

Constraints for Theory Radio spectral features and flares Connection between hard X-ray features and spikes in the range MHz, corresponding to densities of , has always been a promising diagnostic of energy release But there are some aspects hard to understand: frequently the spikes occur in a narrow frequency range for 10s of seconds, implying a fixed density in the energy release site. Energy release widespread over a large volume would produce spikes over a wide frequency (i.e. density) range Wide range of burst types in this frequency range is hard to understand: what controls frequency drift rates of different features?

Radio Emission at Decimetric Wavelengths

Constraints for Theory Magnetic configuration of flares in the low corona See configurations of all types in radio images: single “loops”, double “loops”, complex configurations Frequently see magnetic connections over very large spatial scales Magnetic field strength: spectra typically imply G in the radio source But radio spectra are frequently flat-topped: hard to model, range of fields in the source (need FASR) See both prompt precipitation, implying either rapid scattering of electron pitch angles or loops with little height dependence for B, and trapping, where radio is strong but X-rays are weak, implying little pitch angle scattering.

Radio Flare Loop

The MHD incompressible equations are solved to study magnetic reconnection in a current layer in slab geometry: Periodic boundary conditions along y and z directions Geometry Dimensions of the domain: -l x < x < l x, 0 < y < 2  l y, 0 < z < 2  l z

Description of the simulations: equations and geometry Incompressible, viscous, dimensionless MHD equations: B is the magnetic field, V the plasma velocity and P the kinetic pressure. and are the magnetic and kinetic Reynolds numbers are the magnetic and kinetic Reynolds numbers.

Time evolution of the electric field The surfaces are drawn for E=0.005 from t=200 to t=300

Kinetic energy as a function of time (erg) t (s) t=400 t=300 Total final energy of particles: Magnetic energy:

Energy spectra : e (blue) and p (black) upper panel – neutral, middle – semi-neutral, lower – fully separated beams 1.8 for p 2.2 for e 1.7 for p 4-5 for e 1.5 for p 1.8 for e 4-5 for p 2.0 for e 1.8 for p 2.2 for e

The suggested scheme of proton/electron acceleration and precipitation Pure electron beams, compensated by return current, precipitate in 1s Proton beam compensated by proton-energised electrons precipitate about 10s

Electron Acceleration in Solar Flares basic question: particle acceleration in the solar corona energetic electrons  non-thermal radio and X-ray radiation HXR footpoints HXR looptop electron acceleration mechanisms:  direct electric field acceleration (DC acceleration) (Holman, 1985; Benz, 1987; Litvinenko, 2000; Zaitsev et al., 2000)  stochastic acceleration via wave-particle interaction (Melrose, 1994; Miller et al., 1997)  shock waves (Holman & Pesses,1983; Schlickeiser, 1984; Mann & Claßen, 1995; Mann et al., 2001)  outflow from the reconnection site (termination shock) (Forbes, 1986; Tsuneta & Naito, 1998; Aurass, Vrsnak & Mann, 2002)  radio observations of termination shock signatures

Outflow Shock Signatures During the Impulsive Phase X17.2 flare RHESSI & INTEGRAL data (Gros et al. 2004) termination shock radio signatures start at the time of impulsive HXR rise signatures end when impulsive HXR burst drops off Solar Event of October 28, 2003: The event was able to produce electrons up to 10 MeV.

basic coronal parameters at 150 MHz (  160 Mm for 2 x Newkirk (1961)) (Dulk & McLean, 1978) (flare plasma) shock parameter Discussion I total electron flux through the shock

Summary  The termination shock is able to efficiently generate energetic electrons up to 10 MeV.  Electrons accelerated at the termination shock could be the source of nonthermal hard X- and  -ray radiation in chromospheric footpoints as well as in coronal loop top sources.  The same mechanism also allows to produce energetic protons (< 16 GeV).

Summary The constrains on the acceleration are becoming so many and the ability of a single acceleration to handle all this become impossible- No unique acceleration Shocks, stochastic E-Fields and turbulent acceleration enters into the picture Synchronized from photosheric motions complex magnetic topologies maybe be the answer