Multiwavelength observations of a partially occulted solar flare Laura Bone, John C.Brown, Lyndsay Fletcher.

Outline Background Observations Interpretation Conclusions

Coronal HXR sources First observed in occulted events in the 1970’s using data from OSO 5 and OSO 7 Earliest imaging observations found coronal emission at keV extending to 30000km. Observations with Hinotori extended the energy range to 25keV. Several further sources observed with yohkoh/HXT; – Masuda (1994). – Kosugi et al. (1995). – Feldman et al. (1994). Since the launch of RHESSI in 2002 a number of coronal sources have been observed.

Theoretical interpretations Location of a fast mode shock, occuring where the outflow jet from a coronal reconnection region impacts on a dense and static loop system below. Signature of the current sheet itself Particles trapped and possibly accelerated in the field below a reconnecting coronal structure Thick target bremsstrahlung from non thermal particles in a dense part of the corona.

20 July 2002 X3.3 flare

RHESSI image reconstruction PIXON image reconstruction algorithm used. High quality, excellent noise suppression, photometrically accurate reconstruction. BUT! Very time consuming.

RHESSI images As energy increases, emission concentrated more in looptop, contrary to traditional thick target model.

RHESSI spectra T=26.4MK EM=6.7e49cm -3  =3.9 I(  )=A   =1e6  -3.6

Where the photon spectrum can be approximated by a power law I  A   the instantaneous number of electrons is given by the formula; integrating over energy we can get N(>10keV)=7.0e35 electrons.

OVSA (Owens Valley Solar Array) 2 x 27m and 5 x 2m dishes Tunable to any harmonic of 200MHz from 1-18GHz. Records left, right and circular polarisation.

OVSA data Dynamic spectrum shows impulsive nature of flare From spectrum we can derive different parameters.

Derivation of N and B from Radio data Can fit radio spectrum using a function of the form (Stähli et al.,1989) Where  and  are respectively, the low and high frequency slopes. For the optically thin part of the spectrum shown,   Using; (Dulk and Marsh,1982) we can obtain a value for the electron spectral index of the radio emission to be  =3.13

Assuming a line of sight angle  we can use the polarisation measurements to determine magnetic field strength. r c =0.15 => using the expression given in Dulk (1985); We find at 10.6GHz,  b ~25 at the flare peak (21:30). Thus from; We estimate the magnetic field strength to be ~150G.

We can determine the effective temperature from: and optical depth/electron line of sight density. since S is measured from the radio emission and  estimated from the radio emission, thus we calculate NV=1.5e36 electrons.

Density estimates Can estimate the plasma density in the loop from the emission measure Two separate measurements, RHESSI and GOES. RHESSI EM=6.7x10 49 cm –3 GOES EM=28.0x10 49 cm -3 Gives density estimate of between 1.0x10 11 cm -3 and 2.2x10 11 cm -3. Column density 1-2.2x10 20 cm -2 This implies that electrons of energies 28-41keV being fully stopped in the corona.

Beam driven evaporationc Power >25keV Therefore;

Conductive evaporation Thus for this event; hard for this event to differentiate between beam and conductive evaporation!

Cooling timescales Conductive cooling time assuming constant density and no flows. For values derive gives cooling time of <100sec. Radiative cooling time is given by Which is ~ sec, either conductive cooling is being inhibited or constant heating is occuring.

Conclusions Contrary to typical observations, very high energy electrons observed in a coronal source. ~10 36 electrons instantaneously in the flaring system, V=6.98x10 27 cm 3. Magnetic field ~150G. T~30MK Density >10 11 cm -3,leading to electrons <40keV being stopped in a coronal thick target scenario. Long duration event, mass must be continuously evaporated into the flaring system through both conductive and beam driven evaporation.