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An Overview of Solar Flares with an emphasis on spectroscopy and
the chromosphere
Lyndsay Fletcher University of Glasgow
Hinode 9 Belfast, UK, 15 Sept 2015
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10th September 2014
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Standard model framework
coronal flare loops
Energy flux
Chromospheric footpoint UV/opOcal emission, fast electrons/ions
Post-reconnecOon, field - shrinking and relaxing
ReconnecOon
3
Where & how is flare energy accumulated and stored? Why & how is it released? How is it transported? Where and how is it converted to heat & parOcle KE? How does the atmosphere respond? Energy
flux
evaporaOon
Coronal emission HXR, SXR, EUV, radio
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How does flare energy reach the chromosphere?
Electron (and/or ion) beams from the corona
Thermal conduc7on Alfvenic transport
ParOcle flux
ParOcles lose energy collisionally in the chromosphere
Electron heat flux
Mostly relevant for late phase of flares
reconnecOon reconnecOon
Wave/pulse damps => heaOng and possibly local
parOcle acceleraOon
reconnecOon
Coronal heaOng Coronal acceleraOon
Coronal magneOc disturbance
PoynOng flux
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The flare chromosphere
The chromosphere is of parOcular interest in the study of flares
most flare energy is dissipated here so most flare radiaOon is chromospheric reflects changes in the coronal field offers best diagnosOcs for flare electron distribuOons
Fletcher & Hudson 2001
TRACE 1600A TRACE white light TRACE EUV & HXT
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Flare excess emission
SOL2011-02-15T01:56 excess with EVE/MEGS-A,B,P and Hinode/SOT (Milligan+2014) (5000K blackbody fit to SOT R/G/B) Most of the flare excess radiaOon is emiced in the opOcal, up to a few x 1032 ergs (Woods+2005)
EsOmate from Krucker+2011 flare energy deposiOon rate 1012 ergs cm-2 s-1
For comparison: Solar luminosity equivalent to 6.2 x 1010 ergs cm-2 s-1 Quiet chromospheric heaOng requires ~ 107 ergs cm-2 s-1
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Flare chromosphere in the EUV with EIS
Many spectrocsopic studies of the flare chromosphere with EIS (e.g. Milligan & Dennis 2009, Watanabe+2010, Del Zanna+2010, Graham+2011, Doschek+2013)
Summary of proper7es: Temperature: compact sources cospaOal with ribbons/footpoints observed from log Tmax = 5-7 Density: density at a given T peaks during impulsive phase
ne ~ few x 1010 cm-3 at 2MK (Fe XIV) ne~ few x 1011 cm-3 at 250,000K (O V)
A&A 526, A1 (2011)
Fig. 15. Left: image of the intensity of the blue-wing component inFeXVI (recorded by EIS during the scan No. 6; the field of view is40 120) with two selected regions, a foreground (FG) location inthe core of the active region, and kernel (K) region just a few arcsecondsnorth, in the northern ribbon. Middle: profiles of the FeXVI 263 line,in the two regions. The dashed vertical line indicates the rest velocity,while the dashed lines are the Gaussian profiles of the main and sec-ondary components of the spectrum in the kernel (K). Right: the samespectrum in the K region, fitted with three Gaussian components.
of the main component (dashed line) turns out to be close tothe foreground one, as shown in Fig. 15 (middle). This is asexpected, since we believe that the rest component originatesfrom the foreground active region plasma. We interpret the sec-ondary, blue-wing component as arising solely from the kernelregion, source of the chromospheric evaporation. The intensity,Doppler velocity and width of this component are considerable(see Table 3). An alternative is to consider multiple secondarycomponents having similar widths as the instrumental one. Withtwo such components (bw1 and bw2, see right plot in Fig. 15and Table 3) a good fit is obtained. The two components haveremarkable line-of-sight velocities, about 60 and 170 km s1.
All the strong, un-blended lines formed at similar tempera-tures (in particular the FeXIV 264.787, 274.203 lines) showsimilar characteristics, and were fitted with a two-componentGaussian profile. The strong Fe XV 284.16 line does show thesame features, but it is blended in its blue-wing with Al IX anda recently-identified strong FeXVII line (Del Zanna & Ishikawa2009) which complicates the measurement.
We can rule out the presence of blue-wing enhancementsin lower-temperature lines, because the line profiles of the verystrong FeXIII, FeXII, FeXI, FeX, FeVIII emission lines show noasymmetries. On the other hand, we cannot rule out the presenceof some asymmetry in the hotter lines, because they are eitherweak or blended.
Coronal line profiles with persistent shifts towards the blueor with persistent blue-wing asymmetries have been routinelyobserved by EIS in active region spectra (see, e.g. Del Zanna2008b; Doschek et al. 2008; Hara et al. 2008; Peter 2010).They are often referred to as coronal outflows. We would liketo point out that the blue-wing asymmetries observed here arevery different, for a number of reasons. First, they are very lo-calised at the footpoint of flare loops, contrary to the coronal out-flows, which are also very localised (on plage/Sunspot regions),but then have a very rapid spatial expansion at greater heights(Del Zanna 2008b). Second, they are not observed in FeXIII,FeXII lines, contrary to the coronal outflows case. Third, they arevery transient events, disappearing together with the flare loops.
We can also rule out that the blue-shifts are directly relatedto the ejecta associated with the flare. First, small ejecta appear
Fig. 16. Electron densities (in cm3, in log-scale) from theFeXIV 264.79/274.20 line ratio. The field of view is 38 102 . Thelocation of the ribbons is superimposed.
to lift off from the southern ribbon only, around 23:20 UT, asseen in TRACE 171 . Second, the ejecta are likely to be fila-ment material, having a temperature below 1 MK, the temper-ature of formation of the 171 filter; we have no signaturesof blueshifts at those temperatures in the EIS spectra. The fila-ment ejecta have an obvious importance for the early phase ofthe flare, however higher cadence and sensitivity (as providednow with SDO) would be needed to study this relation.
We now examine in more detail the characteristics of the rib-bon structures. We obtained electron densities from the FeXIV264.79, 274.20 line ratio, using the FeXIV atomic calculationsof Storey et al. (2000). The lines are very strong, hence accuratevalues are obtained. The contribution of the SiVII 274.175 line to the FeXIV 274.203 line has been estimated from thetheoretical intensity ratio with the SiVII 275.383 line (also in-cluded in our EIS study), predicted to be between 0.2 and 0.25in the 1091010 cm3 range. The estimated SiVII contribution tothe rest component of the FeXIV 274.203 line is very small,of the order of 4%. Figure 16 shows the electron densities for allthe EIS rasters.
These were obtained by summing the rest and blue-shiftedcomponents. The average values in the active region are about109.5 cm3, which is typical of an active region (ODwyer et al.2010). However, the ribbon structure becomes denser during thesmall flare, with values in the range 13 1010 cm3. These arehowever lower limits, considering that the densities are averagedvalues of the low-lying ribbons and the foreground plasma.
We have also looked in detail at some kernels. TheFeXIV 264.79, 274.20 line profiles for raster No. 6,23:1623:20 UT, were fitted with two components, one at restwavelength, and the other on the blue-side. Table 3 provides lineintensities, widths and electron densities. Averaged densities of6109 and 11010 cm3 for the rest and blue-wing componentsare found respectively. The density of the rest component, whichcorresponds to the foreground active region emission, is in excel-lent agreement with the density obtained from the FeXIII 202.04,203.83 ratio (5 109 cm3), using the recent new scatteringcalculations of Storey & Zeippen (2010). If one assumes a unityfilling factor (i.e. that the plasma is uniformly distributed alongthe line of sight) and estimates the emission measures for the
A1, page 10 of 14
Lineshi3: chromospheric plasma is strongly blueshi?ed at line formaOon temperatures > log Tmax ~ 6.1 - 6.2 and slightly redshined at lower temperatures Line width: non-thermal broadening (50-60km/s) in footpoints
Fe XIV ne maps, Del Zanna et al. 2010
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Emission measure distribuOon of flare footpoints
EIS EMD is measured in compact flare footpoint sources at impulsive peak of several flares (Graham+2013)
EMD very similar between flares; peaks at log Tmax ~ 6.9, and slope is m ~ 1.
m = 1 consistent with conducOve heaOng of lower chromosphere by a hot upper chromospheric source, balanced by radiaOve cooling (Shmeleva & Syrovatskii 1973)
T ~ 10 MK Fcond
Qrad
m = 1
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Chromospheric ribbons in the UV with IRIS
Li & Zhang (2
015)
80
Endpoints of just-reconnected field => informaOon about topological changes & reconnecOon
rapid movement of bright sources along ribbons
ribbon front someOmes fluctuates/reverses
Endpoints are not pointlike.
bright leading edge, scale ~1 fibril structure in places 1400 SJC
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Low-temperature response to energy input on a fine scale
Li & Zhang (2
015)
Li & Zhang (2015) Single gaussian fisng of Si IV 1403 (80,000K) sit-and-stare data: Repeated bright fluctuaOons that run along ribbon cross into the slit Each pulse is associated with transient net redshin and increased net linewidth
intensity
centroid
Line width
Origin of broadening.Unresolved flows? Turbulence? Mul7-thermal DEM? Non-Maxwellian distribu7ons? Pressure/opacity effects?
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28
Si IV 1403 intensity
-190
-188
-186
-184
-182
Sola
r-Y (a
rcse
c)
12:45 12:50 12:55 13:00 13:052014 Apr 18
Si IV 1403 intensity
800 1000 1200 1400 1600 1800Time (s) since 12:33:38 UT
-190
-188
-186
-184
-182
Sola
r-Y (a
rcse
c)
1 2 3 4 5 6
3.0
3.5
4.0
4.5
5.0log10(Total DN/s)
Fig. 6. Upper panel: time-distance stackplot of the total Si iv 1403 A SG passbandintensity, in RLBW. Time is given on the x-axis in UT, and the y-axis is solar-Y in arcsec.
Intensity scale is given to the right. Lower panel: reprint of the upper panel with a redoutline indicating the position of the sawtooth pattern described in Section 3.1. Time isgiven on the x-axis in s (after 12:33:38 UT), and the y-axis is unchanged. The blue line
indicates the sawtooth centroid position, and the orange line is a linear fit to the blue line.The numbers 1-6 indicate six peaks in the sawtooth oscillation.
Fisng with a double Gaussian
SOL2014-04-18T13:03 Brannon+2015 fit Si IV with a double gaussian. One component oscillates between blue and redshin, and another which is more consistently red-shined
Si IV Intensity
Doppler speed
29
-50 0 50 100 150v [km/sec]
A
12:5
1:57
-50 0 50 100 150v [km/sec]
-50 0 50 100 150v [km/sec]
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12:5
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-50 0 50 100 150v [km/sec]
12:5
6:01
-50 0 50 100 150v [km/sec]
B
12:5
0:52
12:5
1:20
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6:30
12:5
6:49
12:5
7:26
12:5
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C
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0:14
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1:39
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4:00
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D12
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52
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-50 0 50 100 150v [km/sec]
E
12:5
3:41
-50 0 50 100 150v [km/sec]
-50 0 50 100 150v [km/sec]
12:5
4:00
-50 0 50 100 150v [km/sec]
-50 0 50 100 150v [km/sec]
12:5
4:18
-50 0 50 100 150v [km/sec]
-50 0 50 100 150v [km/sec]
12:5
4:46
-50 0 50 100 150v [km/sec]
A
B C
D E
Fig. 7. Selected Si iv spectral lines for 28 pixels in or near the sawtooth pattern shown
in Figure 6. An outline of the sawtooth is shown at upper right, and the five rows of plotsA-E each sample a given spatial position at several different times, as indicated by the+ symbols at upper right. Each plot displays the spectrograph data for Si iv 1394 A (as
asterisks) and 1403 A (as squares), with the time printed to the right of each plot. Thex-axis is in km s1 from nominal line center (same for both lines), and the y-axis scaling is
arbitrary (the 1394 A data has been scaled by a factor of 0.5 for direct comparison to the1403 A line). Additionally, the solid orange (green) lines are the two component Gaussianfits for the 1394 A (1403 A) spectral line, as described in Section 3.2. The dashed lines are
the individual components CB and CR for each of the Si iv spectral lines, with cyan/violetused for the respective 1394 A spectral components and blue/red used for the respective
1403 A spectral components.
34
LOS
North
NorthLOS
Sub-chromospheric anchor
Ribbon motion
Blueshift
Redshift
Flaring loop
Chromosphere
Chromosphere
Elliptical loop motion
Current sheet
KH or TM instability
Reconnection
171 loop(~30 min)
Fig. 12. Upper panel: schematic diagram of our proposed scenario, with a KH or TMinstability in the coronal current sheet resulting in elliptical wave oscillations in the recon-
nected flare loops. Lower panel: schematic diagram of the flare loop footpoint, as describedin Section 4, showing how the elliptical wave motion relative to the LOS generates the
observed 180 phase difference between the ribbon motion and Doppler velocities.
InterpretaOon: - ellipOcally polarised oscillaOon of field in transiOon region
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Chromospheric upflows
With the Fe XXI line and chromospheric lines we can examine chromospheric evaporaOon, condensaOon and momentum balance in detail
Some pacerns emerging from various studies, reinforcing EIS results. Upflows of 100s of km/s in Fe XXI
Licle/no staOonary hot plasma at footpoints
SubstanOal line broadening at footpoints
0 km/s = Fe XXI rest posiOon km/s
Youn
g et al. 2015
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Flare opOcal spectrum
The Astrophysical Journal, 798:107 (17pp), 2015 January 10 Kowalski, Cauzzi, & Fletcher
2x105 0 2x105 4x1050
10
20
30
2x105 0 2x105 4x1050
10
20
303670 3
Umbral Kernels
2x105 0 2x105 4x1050
10203040 3920
2x105 0 2x105 4x10505
1015202530 4170
2x105 0 2x105 4x1050
10203040 4450
2x105 0 2x105 4x105 6x105 8x105 1x106
Excess Intensity
01020304050
H
Figure 7. Histograms of the excess intensity (subtracting 15:07:24 from15:09:30) over a spatial cut through the HSG data, at selected continuumintervals and H . The value corresponding to 3 of the distribution is indicatedby dashed lines, whereas the values for the 3 pixels averaged to obtain the excessI for the umbral kernel are indicated by dotted lines. Note that the value of in this figure represents the spatial variation of the excess, whereas the valueof in Figure 6 describes the temporal variation of the excess.
4.1. Comparison with the Hard X-Ray Fermi Light Curve
In Figures 10(a)(e) we show the continuum-subtracted, line-integrated excess intensity as a function of time in H and Ca iiK for the first slit position at the location of the umbral kernel(a), and in adjacent regions to the umbral kernel in the second,third, fourth, and fifth slit positions ((b)(e), respectively).The uncertainties of the integrated excess line intensity werecalculated following the standard formula in the Appendix ofChalabaev & Maillard (1983), which adds the uncertainty of theintegrated continuum and line excesses in quadrature. Althoughthe emission line excess extends over a larger region (3;Figure 4) compared to the continuum excess, we average theintensity only over the three brightest pixels along the slit. Theemission line light curves are compared to the 1521 keV hardX-ray light curve obtained from Fermi/GBM and the 18 soft X-ray luminosity obtained from GOES. The same graybars from Figure 6 indicate the times of significant continuumexcess in panel (a). In the leftmost slit position of the raster (a),we see a general similarity in the normalized time variation ofthe hard X-rays and the optical lines. As we progress away fromthe leftmost slit position (panels (b)(e)), we observe a moregradual response in the optical lines, yet reaching a comparablemaximum value as in the first slit position. By the fifth slitposition, i.e., 2.5 apart, the optical lines appear to evolvesimilarly to the soft X-ray emission. The gradual evolution inthe chromospheric emission lines coincides with the formationof new, low-brightness kernels at these later times, as seen inthe development of the H+1.2 umbral ribbon (Figure 5).
At the location of the umbral continuum excess, we findcoincident peaks between the hard X-ray and the opticalemission line light curves, but the 20 s cadence of the opticallight curves makes it difficult to compare the precise timing withthe much better sampled hard X-ray light curve (t 4 s). Thefirst enhancement at 15:08:27 in the hard X-rays correspondswell to the first impulsive enhancement in the optical lines, butwe do not observe a significant continuum peak in Figure 6at this time. The maximum of the hard X-rays at 15:09:25corresponds to a major peak in both optical lines (15:09:30) and
3600 3800 4000 4200 4400
0
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3600 3800 4000 4200 4400Wavelength ()
0
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Exc
ess
I (
105
erg
s1 c
m2 s
r1
1)
BJ
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H8
Ca
II K
Ca
II H
+H
H
H
15:09:30
0.00
0.05
0.10
0.15
0.20
Fla
re C
ontr
ast
Excess I/Io
Figure 8. Full spectral range of the excess intensity at 15:09:30 in the umbral kernel at the first slit position in the spectral raster. The intensity is averaged over threespatial pixels. The right axis (square symbols) show the flare contrast in selected continuum wavelength regions.
9
Broad-band flare observaOons in the opOcal are very rare only a few spectra from 70s & 80s (Neidig 83, DonaO-Falchi+84)
Some show evidence for Balmer recombinaOon edge (i.e., free-bound emission) others do not.
SOL2011-08-18 observed at the Dunn Solar Telescope by the Horizontal Spectrograph - high dispersion spectrum from 3600-4550 (Kowalski+14)
The Astrophysical Journal, 798:107 (17pp), 2015 January 10 Kowalski, Cauzzi, & Fletcher
0 10 20 30 40 500
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20
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Broadband 15:06:58
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Halpha + 1.2 15:09:34
0 10 20 30 40 500
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Halpha + 1.2 15:11:53
Figure 1. Central portion of the IBIS field of view at various times during flare development. Axes are in units of arcseconds. Top left: pre-flare broadband image at6360 . The two vertical black lines outline the edges of the HSG raster described in Section 2.3 and displayed in Figure 4. Top right: H+1.2 , at the same preflaretime. An early brightening along the (future) flare ribbon is already visible within the larger spot. Bottom left: H+1.2 near the time of largest hard X-ray peak inthe Fermi curve (see Figure 2 and Figure 10). The left white line indicates the HSG slit position at 15:09:30 and the right white line indicates the HSG slit positionat 15:08:44. Excess continuum was detected in the small flare kernel crossed by the slit around position (26, 39) at this time. The right white line indicates theHSG slit position crossing the plage flare ribbon described in Section 4.2. Bottom right: H+1.2 at a later time during the flare development. Note the motion of theplage flare ribbon away from the earlier position. All H images are scaled to the same intensity levels. Note that images maintain the native orientation, with verticaldirection along the parallactic angle; the east limb direction is roughly toward the bottom of the figure and north is to the right.
15:06 15:08 15:10 15:12 15:14Time (UT) 18Aug2011
0.00
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mi 1
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ux (
cts
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ES
Xr
ay L
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24 e
rg s
1)
Figure 2. Fermi 1521 keV hard X-ray count flux light curve (left axis) shown with the GOES 18 luminosity (right axis) of the C1.1 flare SOL2011-08-18T15:15from AR 11271. The timing of the simultaneous H + 1.2 and optical continuum enhancements are indicated by vertical gray bars. Each Fermi data point has a livetime of 4.07 s.
3
HSG slit posiOon at flare peak
White light kernel
Flare excess spectrum no clear Balmer jump. Though could be hidden by merged, broadened
high-order H lines
BJ
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Flare opOcal conOnuum
components gives consistent results (see Dennis & Per-nak 2009). The emission perpendicular to the ribbon isunresolved at RHESSIs resolution of 2. 3 (FWHM). TheHXR footpoint is seen down to 20 keV, with the 2030 keVimage showing the same source morphology as at higherenergies, but with a centroid being insignificantly higher by33 km. At even lower energies, the thermal emission from thecoronal flare loop dominates and emission from the footpointsource is lost in the limited dynamic range of the RHESSIimages. The time evolution of the HXR emission revealssimilar altitudes for the other 10 time intervals during theimpulsive phase (Figure 2), with a trend toward slightly loweraltitudes with time (about 300 km over the 7 minutes of themain HXR burst) that will be discussed below. The averageddifference of the WL and HXR altitude is 31 km with astandard deviation of 98 km. That the standard deviation issimilar to the estimated statistical errors indicates that the errorestimates are trustworthy. Within the statistical error, weconclude therefore that the HXR and WL are co-spatial.
A trend toward lower altitudes in time is observed for WLand HXR centroids (Figure 2, left bottom). This trend is mostlikely due to an apparent motion when the source positionschange, but could also be due to an actual height change (seeFigure 1). Besides this apparent motion in the radial direction,the northern source is also seen to systematically move alongthe limb in time with an average projected speed of 8 km s1(Figure 7). This apparent footpoint motion is generally thought
to be due to the geometrical development of the reconnectionprocess (e.g., Krucker et al. 2005). The northern ribbonstructure seen in STEREO suggests that the ribbon directiondeviates only slightly from the line of slight as seen from Earth.Nevertheless, it is unclear what the observed motion along thelimb corresponds to. It could be due to separation of the ribbonas well as a foreshortened motion of a motion along the ribbon,or a combination of these two. Motions along the ribbons couldproduce the observed trend toward lower altitudes in time. Theobserved decrease by 300 km could be the result of a change inoccultation heights by that amount. Starting exactly at the limb,this would correspond to a change by 2. 9. This corresponds toa distance of about 40Mm, and considering the 7 minuteduration, a projected footpoint speed of 100 km s-1. This is arapid footpoint motion speed, but not untypical for large flares(e.g., Metcalf et al. 2003; Yang et al. 2009). On the other hand,the estimated occultation height could be lower as the motioncould start not right at the limb, and then a smaller angle couldgive the same effect. Despite all these unknowns, it is quiteplausible that the decreasing altitude is an effect of footpointmotion. In such a case, the highest altitudes observed give thebest estimates of the actual altitude with values around 900 km.Figure 7 further reveals that the HXR sources outline the
leading edge of the WL source, consistent with the picture thatthe energization of the WL emission is due to the non-thermalelectrons that produce the HXR signal. The trail of the WLemission behind the newly flaring part of the ribbon could
Figure 5. Zoomed view (15 15) around the stronger flare ribbon shown in Figure 3: The top row shows RHESSI Clean images reconstructed with the nominalhighest resolution of 2 . 3 FWHM. The RHESSI images shown below are reconstructed with an artificially reduced point-spread function of 1.1 FWHM to roughlymatch the resolution of the HMI images. The same contours are shown for the RHESSI (blue) and HMI (red) images at 50, 70, 90% (for the more complex flare ribbonof the November 20 flare additional levels at 10 and 30% are shown).
6
The Astrophysical Journal, 802:19 (10pp), 2015 March 20 Krucker et al.
Krucker+2015
HMI WL 79970 km RHESSI 30-80keV 74651 km
What makes the flare UV-opOcal conOnuum? Free-bound emission? Enhanced photospheric blackbody? Direct heaOng? backwarming? Whatever the mechanism, the energy must reach deep into the atmosphere
OpOcal and HXR bremsstrahlung sources are co-spaOal
Formed beyond expected collisional range of HXR-emisng electrons
Perhaps chromosphere is underdense compared to expectaOons? Perhaps electrons are accelerated in situ in chromosphere? RadiaOon hydro models may shed light on this
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Pusng it together
Transient high temperature plasma, up to 10MK, in upper chromosphere (where ne < 1012 cm-3)
EMD suggests possibility of conducOve heaOng down to log T = 5 (But also accessible by electron beams) Plasma flows along the field (evaporaOon/condensaOon) but also
possible fluctuaOons of the field itself
Broadening of hot lines may suggest plasma turbulence Lower chromosphere produces opOcal emission modest
temperature rises but energeOcally dominant
Not at all clear how to get energy to the lower chromosphere
Summary
The research leading to these results has received funding from the European Communitys Seventh Framework Programme (FP7/2007-2013) under grant agreement no. 606862 (F-CHROMA)