OBSERVATIONS OF THE THERMAL AND DYNAMIC EVOLUTION OF A SOLAR MICROFLARE

Transcription

1 The Astrophysical Journal, 692: , 2009 February 10 c The American Astronomical Society. All rights reserved. Printed in the U.S.A. doi: / x/692/1/492 OBSERVATIONS OF THE THERMAL AND DYNAMIC EVOLUTION OF A SOLAR MICROFLARE Jeffrey W. Brosius 1 and Gordon D. Holman 2 1 Catholic University of America at NASA Goddard Space Flight Center, Solar Physics Laboratory, Code 671, Greenbelt, MD 20771, USA; Jeffrey.W.Brosius@nasa.gov 2 NASA Goddard Space Flight Center, Solar Physics Laboratory, Code 671, Greenbelt, MD 20771, USA; Gordon.D.Holman@nasa.gov Received 2008 August 14; accepted 2008 October 24; published 2009 February 19 ABSTRACT We observed a solar microflare over a wide temperature range with three instruments aboard the SOHO spacecraft (Coronal Diagnostic Spectrometer (CDS), Extreme-ultraviolet Imaging Telescope (EIT), and Michelson Doppler Imager (MDI)), TRACE (1600 Å), GOES, and RHESSI. The microflare s properties and behavior are those of a miniature flare undergoing gentle chromospheric evaporation, likely driven by nonthermal electrons. Extremeultraviolet spectra were obtained at a rapid cadence (9.8 s) with CDS in stare mode that included emission lines originating from the chromosphere (temperature of formation T m K) and transition region (TR), to coronal and flare (T m K) temperatures. Light curves derived from the CDS spectra and TRACE images (obtained with a variable cadence 34 s) reveal two precursor brightenings before the microflare. After the precursors, chromospheric and TR emission are the first to increase, consistent with energy deposition by nonthermal electrons. The initial slow rise is followed by a brief (20 s) impulsive EUV burst in the chromospheric and TR lines, during which the coronal and hot flare emission gradually begin to increase. Relative Doppler velocities measured with CDS are directed upward with maximum values 20 km s 1 during the second precursor and shortly before the impulsive peak, indicating gentle chromospheric evaporation. Electron densities derived from an O iv line intensity ratio (T m K) increased from cm 3 during quiescent times to cm 3 at the impulsive peak. The X-ray emission observed by RHESSI peaked after the impulsive peak at chromospheric and TR temperatures and revealed no evidence of emission from nonthermal electrons. Spectral fits to the RHESSI data indicate a maximum temperature of 13 MK, consistent with a slightly lower temperature deduced from the GOES data. Magnetograms from MDI show that the microflare occurred in and around a growing island of negative magnetic polarity embedded in a large area of positive magnetic field. The microflare was compact, covering an area of km 2 in the EIT image at 195 Å, and appearing as a point source located 7 west of the EIT source in the RHESSI image. TRACE images suggest that the microflare filled small loops. Key words: Sun: activity Sun: corona Sun: flares Sun: transition region Sun: UV radiation Sun: X-rays, gamma rays 1. INTRODUCTION It has long been known that solar active regions exhibit impulsive EUV bursts that suggest the occurrence of microflares (Lites & Hansen 1977; Emslie & Noyes 1978; Porter et al. 1984). Using the ultraviolet spectrometer aboard Orbiting Solar Observatory 8 (OSO-8) in an operating mode that scanned spectral line profiles, Bruner & Lites (1979) scanned successive profiles of the C iv 1548 Å resonance line (formed at temperature T m K) at intervals between 27 and 50 s, and found that redshifts accompanied most of the impulsive brightenings. Athay et al. (1980) concurred, but pointed out that in the presence of a rapidly changing intensity (some bursts vary on timescales shorter than the minimum profile scan time) the line profile centroid is artificially altered by the time delay in scanning the two sides of the line profile. (Throughout this work we refer to an emission line s formation temperature T m as the temperature at which the given ion s ionization fraction is maximized (hence the subscript m ); see Mazzotta et al ) Emslie & Noyes (1978) examined light curves of EUV impulsive bursts observed at seven different wavelengths with the EUV spectroheliometer aboard Skylab, and found that for each burst all seven light curves were synchronized within the 5.5 s time resolution of the observations. They concluded that a low flux of nonthermal electrons was a suitable candidate for the burst energy source. Lin et al. (1984) used a highly sensitive balloon-borne instrument to detect hard X-ray bursts that lasted from a few 492 seconds to several tens of seconds. Termed microflares, the hard X-ray bursts were generally accompanied by small soft X-ray bursts. Calculations based on single temperature thermal sources were insufficient to fit the observed spectra; power-law fits were required. This suggested that even very small transient releases of energy by the Sun may be primarily nonthermal in character. Porter et al. (1995) report observations of an active region with the UltraViolet Spectrometer/Polarimeter (UVSP) and Hard X-ray Imaging Spectrometer (HXIS) aboard the Solar Maximum Mission (SMM) during a period in which there were no major flares. Because many events substantially smaller than subflares (identified in C iv 1548 Å emission) were found to have impulsive counterparts in kev X-ray emission characteristic of K plasma, these events were also termed microflares and considered to be true members of the flare family. Krucker et al. (2002), Benz & Grigis (2002), and Stoiser et al. (2007, 2008) used the Ramaty High Energy Solar Spectroscopic Imager (RHESSI) (Lin et al. 2002; Hurford et al. 2002; Smith et al. 2002) satellite to investigate hard X-ray microflares down to 3 kev. Some events showed elongated sources while others were unresolved point sources. During the impulsive phase, the microflare spectra were best fitted with a thermal component for energies below 7 10 kev, and a nonthermal power law at higher energies. Jain et al. (2006) found a similar result based on data from the SOlar X-ray Spectrometer (SOXS; Jain et al. 2005). Stoiser et al.

2 No. 1, 2009 THERMAL AND DYNAMIC EVOLUTION OF A SOLAR MICROFLARE 493 (2007, 2008) also obtained coordinated 1600 Å and 171 Å emission with the Transition Region And Coronal Explorer (TRACE) (Handy et al. 1999) satellite, which light curves peaked simultaneously with the RHESSI light curves (within the 1 minute time resolution of the TRACE images). This suggests that the standard flare model in which precipitating electrons heat cold, chromospheric material to flare temperatures also applies to microflares. This process, first described by Neupert (1968), later became known as chromospheric evaporation (Antonucci et al and references therein; Bornmann 1999 and references therein; Czaykowska et al. 1999, 2001; Brosius2003; Teriaca et al. 2003, 2006; Allred et al. 2005; Falchi et al. 2006). When available, microwave observations of microflares corroborate the nonthermal character of the hard X-ray emission (Kundu et al. 2006; Qiu et al. 2004). The comparison by Stoiser et al. (2008) of model calculations (involving electron beam and conductively driven chromospheric evaporation in single and filamented flare loops) with the same microflares presented by Stoiser et al. (2007) supports beam heating in filamented loops as the most likely scenario. Chromospheric evaporation occurs when chromospheric material is heated by beamed particles and/or thermal conduction more quickly than it can radiatively cool. The heated material expands upward into the lower-density corona, and slowly downward into the higher-density chromosphere. Evaporation is said to proceed gently when emission lines formed at temperatures characteristic of the upper chromosphere and transition region (TR) (as well as hotter lines, if observed) all appear blueshifted; evaporation is said to proceed explosively when emission lines formed at temperatures characteristic of the upper chromosphere and TR all appear redshifted while hotter lines appear blueshifted (Fisher et al. 1985b; Brosius & Phillips 2004; Milligan et al. 2006a, 2006b; Brosius & Holman 2007). In their hydrodynamic simulations of the response of a flare loop atmosphere to thick-target electron heating, Fisher et al. (1985a) found that a beam energy flux of about erg cm 2 s 1 served as an effective threshold between gentle and explosive evaporation. Christe et al. (2008) conducted a statistical survey of 25,705 RHESSI microflares (below GOES C-class) identified with an automated flare-finding algorithm. All of the microflares occurred in active regions. Hannah et al. (2008) conducted a statistical analysis of the thermal and nonthermal properties of the microflares identified by Christe et al. (2008). Each microflare was analyzed automatically at the peak time of its 6 12 kev emission, with reliable thermal plus nonthermal fits derived for 4236 of the microflares identified. Median values of several microflare parameters derived from the survey include a duration of 5.4 minutes, a thermal source temperature of 12.6 MK, an emission measure of cm 3, and a thermal loop length of 32. To obtain rapid cadence (9.8 s) EUV flare spectra with the Coronal Diagnostic Spectrometer (CDS; Harrison et al. 1995) aboard the Solar and Heliospheric Observatory (SOHO) spacecraft, the observing sequence PBEAM (Brosius 2001, 2003) was developed. This is achieved by (1) reducing the spatial resolution along the slit by rebinning (averaging) the spatially resolved spectra into twelve 4 20 pixels; (2) recording only Å and Å subsets of the CDS Å waveband; and (3) staring at (not rastering over) the target. Brosius & Phillips (2004) presented PBEAM observations of a GOES M2.3 flare loop footpoint. Based on upflows 40 km s 1 measured in the relatively cool EUV lines of O iii, Oiv, Ov, and He ii during two flare precursors (with no simultaneous downflows), they concluded that gentle chromospheric evaporation occurred at those times. The same lines exhibited downflows 40 km s 1 later during the flare impulsive phase, when the hot flare line of Fe xix became strong and blueshifted, with a maximum upflow speed of 64 km s 1. This, along with comparable upflows measured in hot flare lines of S xv (T m K) and Ca xix (T m K) observed with the Yohkoh satellite s Bragg Crystal Spectrometer, indicated explosive evaporation. The CDS observing sequence FLAREDOP is a modified version of PBEAM that obtains spectra at wavelengths of , , and Å. Emission lines in these spectra cover a wider range of formation temperature than that available to PBEAM, including He i at and Å (T m K), He ii Å (seen in second order, T m K), O iii Å (T m K), O iv and Å (T m K), O v Å (T m K), Mg x at and Å (T m K), Si xii at Å (T m K), and Fe xix Å (T m K). Brosius & Holman (2007) presented FLAREDOP and coordinated RHESSI observations of a small flare-like transient (FLT) observed during a GOES M1.6 solar flare at a remote location 1 arcmin or more away from the flare itself. EUV light curves derived from CDS spectra revealed a brief precursor and an impulsive peak followed by a more gradual rise and decline of emission; hard X-ray light curves obtained with RHESSI revealed a small burst just before the EUV impulsive rise, and another burst at the time of the more gradual EUV peak. During the impulsive phase, simultaneous downward velocities 30 km s 1 were measured in the chromospheric line of He i at Å and the TR line of O v at Å, along with upward velocities 20 km s 1 in the coronal line of Si xii at Å. These oppositely directed Doppler velocities, along with the emergence of weak, blueshifted Fe xix emission at Å during the impulsive phase, demonstrated that explosive chromospheric evaporation occurred at this location remote from the primary region of particle acceleration. The present work is based on CDS FLAREDOP spectra of a microflare that was observed simultaneously with the GOES, RHESSI, and TRACE satellites, as well as with the Extreme-ultraviolet Imaging Telescope (EIT; Delaboudinière et al. 1995) and Michelson Doppler Imager (MDI; Scherrer et al. 1995) aboard SOHO. We were fortunate to have had the CDS slit point directly at the compact microflare before, during, and after the event. Although RHESSI detected no photons at energies above 10 kev and found no evidence of nonthermal emission, the microflare s behavior observed in EUV emission was very similar to that of a larger solar flare undergoing gentle chromospheric evaporation. Its source was located in and around a growing area of negative magnetic polarity embedded in a larger area of positive magnetic field. In Section 2, we describe the observations and data reduction procedures, in Section 3 we present the results of our analysis, in Section 4 we discuss implications of our results, and in Section 5 we summarize our conclusions. 2. OBSERVATIONS AND DATA REDUCTION We observed a GOES B2 solar microflare (that peaked at the A5 level after background subtraction) in NOAA AR on 2005 November 16 around 08:12 UT with RHESSI, TRACE, and SOHO s EIT, CDS, and MDI. There was only one numbered

3 494 BROSIUS & HOLMAN Vol. 692 Figure EIT 195 Å images obtained at (a) 07:59:52 UT and (b) 08:11:51 UT on 2005 November 16, both displayed on the same negative intensity scale. The microflare appears only in frame (b), where it is indicated with an arrow and highlighted with a contour that corresponds to 75% of the maximum EIT intensity in the microflare. The position of the CDS slit with its twelve 4 20 pixels is outlined in red. region on the disk at the time, and the Sun was very quiet, with just one GOES flare (a C5.9 event at 11:41 UT) reported on this date EIT Under quiescent conditions the EIT 195 Å waveband is dominated by Fe xii line emission (192.4, 193.5, Å) formed at T m K. During flares, however, Fe xxiv line emission (192.0 Å) formed at T m K may contribute. Based on the usual synoptic sequence of 195 Å full-disk images obtained at 12 minutes cadence with EIT, our microflare is evident only in the image obtained at 08:11:51 UT. See Figures 1 and 2(a). Compared to the next earliest EIT image (07:59:52 UT) the brightening encompassed 11 contiguous pixels with enhancement factors that ranged from 1.40 to 2.29, with an average value of The corresponding solar surface area was about 75 arcsec 2,or km 2. The brightening was gone before the start of the next 195 Å image obtained at 08:23:51 UT. We found that the 195 Å image center coordinates listed in the data file headers were incorrect by more than three pixels in the y-direction and less than one pixel in the x-direction. We derived a corrected set of center coordinates for each image using the IDL procedure eit_point, and verified them by visual inspection before using them throughout this work. The largest fractional intensity increase in the 195 Å image at 08:11:51 UT relative to the corresponding pixel intensities in the 07:59:52 UT image was centered at ( , ); this position also corresponds to the brightest microflare pixel in the EIT image. Here and throughout this paper we give positions in solar disk coordinates as viewed from 1 AU (not as viewed from SOHO s position at L1) MDI Based on the usual synoptic sequence of full-disk photospheric longitudinal magnetograms obtained at 96 minutes cadence with MDI, along with a few magnetograms obtained at 1 minute cadence that are available beginning about 40 minutes after the microflare, the microflare is associated with a growing area of negative magnetic field that is surrounded by a large area of positive magnetic field toward the center of the active region. Figure 2(b) shows a portion of the magnetogram obtained at 08:03:02 UT; this is the magnetogram closest in time to the microflare. Black areas represent inward-directed field, and white areas represent outward. The green X indicates the maximum intensity observed in the EIT 195 Å image obtained at 08:11:51 UT, and the blue X indicates the maximum count rate observed in RHESSI s 3 10 kev energy bin integrated between 08:12:08 and 08:12:40 UT (see below). Red lines outline the 4 20 pixels in which the CDS spectra were obtained TRACE TRACE observed at a variable but relatively rapid cadence (30 55 s, with an average value of 34 s) in its 1600 Å passband during the event. This 275 Å wide passband includes lines from C i and Fe ii as well as continuum emission, so that the temperature range covered extends from about 4000 K to 10,000 K. The TRACE observations represent the coolest plasma in our microflare dataset. The actual TRACE pointing may be offset up to 10 from the nominal pointing listed in the data file headers (e.g., Metcalf 2001), so we determined the actual pointing for our observations by coaligning a TRACE continuum image (obtained occasionally during our observing sequence) with a full-disk continuum image from MDI obtained at the same time (08:37 UT in this case). Excellent coalignment of the same solar features observed with the two instruments demonstrates that the actual center of the TRACE field of view (FOV) was located 4 westward and 6 southward of the value listed in the header. Applying this correction (here and throughout this work), we find that the microflare appears in 1600 Å imagery near the same location as the EIT (and RHESSI) source. See Figure 2(c). A sequence of TRACE images, shown in Figure 3, reveals a compact source that evolves rapidly in time. The FOV in this figure is the same as in Figure 2(c). Some of the structures give the appearance of short, low-lying loops that fill and brighten as the event proceeds. We obtained a light curve for the TRACE 1600 Å passband by averaging its intensities over the 8 40 array of 0. 5 pixels that corresponds to the 4 20 CDS slit pixel in which the microflare was observed. The resulting light curve tracks those of the photospheric and TR emission lines observed with CDS and displayed in Figure 4. However, because the TRACE time resolution was only about one-third that of CDS during these observations, we do not include the 1600 Å light curve in this figure CDS The CDS FLAREDOP study was used to obtain a complete set of , , and Å EUV spectra in each of the twelve 4 20 slit spatial pixels every 9.8 s

4 No. 1, 2009 THERMAL AND DYNAMIC EVOLUTION OF A SOLAR MICROFLARE 495 Figure 2. Four views of the microflare and its environs, including (a) an EIT 195 Å image obtained at 08:11:51 UT, (b) an MDI photospheric longitudinal magnetogram obtained at 08:03:02 UT, (c) a TRACE 1600 Å image obtained at 08:11:54 UT, and (d) a RHESSI 3 10 kev pixon image integrated between 08:12:08 and 08:12:40 UT. The FOV in each frame is slightly different because each frame displays a whole number of pixels for its respective instrument, most of which are not integral multiples of each other. Red lines outline the 4 20 CDS slit pixels, and the microflare is indicated with an arrow. In frame (a) the contours correspond to 75% and 90% of the microflare s maximum 195 Å intensity measured with EIT; the white X indicates the location of the RHESSI source. In frame (b) white represents outward-directed (positive) magnetic field and black represents inward; the green X indicates the location of the maximum 195 Å intensity measured with EIT, and the blue X indicates the location of the RHESSI source. In frame (c) the contours correspond to 15% and 50% of the microflare s maximum 1600 Å intensity measured with TRACE; X s again indicate the maxima measured with EIT and RHESSI. In frame (d) the contours correspond to 15% and 50% of the microflare s maximum 3 10 kev emission measured with RHESSI; the green X is as described above. while the instrument stared at the active region. The spectra were processed and calibrated with standard SolarSoftware IDL procedures. We used the broadened Gaussian line profile fitting procedure developed by Thompson (1999b) to obtain profile fits to emission lines. Brosius (2003) and Brosius & Phillips (2004) discuss why and how this procedure is used to obtain centroid wavelengths and integrated line intensities. In what follows we rely most heavily on the strong, unblended lines of He i at Å, O v at Å, and Si xii at Å. The Fe xix line at Å is weak and detectable for only about 2 minutes during the flare. Light curves for these four lines are shown in Figure 4. It is well known that the actual pointing of CDS may differ by about 10 from the commanded value (Thompson 1999a). Therefore, in order to derive the most accurate possible position of the CDS slit for the stare observations analyzed here, we used the 195 Å intensities observed by EIT (usually dominated by Fe xii line emission formed at T m K) to fine tune the CDS slit pointing. This was done by averaging selected EIT pixel intensities into spatial bins appropriate for those into which the CDS spectra were recorded, and adjusting the origin of those bins until the spatial dependence of the EIT 195 Å intensity matched that observed along the CDS slit in lines formed at comparable temperatures. Since Mg x observed with CDS has nearly the same formation temperature as Fe xii, we used Mg x lines to achieve the coalignment. This application of the Mg x lines (609.8 and Å) is not compromised by their known blending with O iv lines (see Brosius & Phillips 2004) since Mg x dominates the blends especially during quiescent times. We used the nominal position of the CDS slit as the starting point for the EIT pixel binning, and adjusted the slit s position within a range limited by the known uncertainty in the CDS pointing until the scatter in the ratio of CDS to EIT pixel intensities was minimized. The process was repeated for the four EIT images obtained at 07:47:51, 07:59:52, 08:11:51, and 08:23:51 UT, using Mg x intensities from CDS slit spectra at the same times. This yielded a shift (relative to header value) of to the east and 2. 3 to the north, for a net shift of The position of the slit thus derived is shown in Figures 1, 2, and 3. We investigated relative Doppler velocities during the microflare by defining rest wavelengths for He i at Å, O v at Å, and Si xii at Å to be average values of

5 496 BROSIUS & HOLMAN Vol. 692 Figure 3. Sequence of TRACE 1600 Å images of the microflare covering the same ( ) FOV displayed in Figure 2(c). Time (UT) is given in the upper left of each frame. Red lines outline the 4 20 CDS slit pixels. Figure 4. Light curves derived from CDS spectra of EUV emission lines formed over a wide range of temperature, from the chromosphere (He i Å) through the TR (O v Å) and into the corona (Si xii Å), including hot flare emission (Fe xix Å). The color scheme is indicated in the upper left of the figure. For clarity of display, the He i line intensity was reduced by a factor of 1.5 while the Si xii and Fe xix line intensities were multiplied by factors of 7 and 10, respectively. The microflare is preceded by two precursor brightenings (07:58:16 08:00:42 UT, indicated with dotted vertical lines, and 08:01:51 08:08:23 UT, indicated with dashed vertical lines) before itself begins about 08:09:12 UT and ends about 08:16:03 UT (solid vertical lines) as described in the text. The microflare is evident in its chromospheric and TR emission for 2.5 minutes before its coronal emission begins to increase, and 2.8 minutes before its hot flare emission appears above the noise. centroid wavelengths derived from line profile fits in rapid cadence spectra obtained between 07:30:01 and 07:56:57 UT before the microflare (and before its precursors), all within the CDS spatial pixel that actually observed the microflare. The associated 1σ scatters were also derived (and converted to velocity units) in order to assess the significance of the velocities derived during the microflare. For the above lines of He i,ov, and Si xii we obtain 1σ values of 1.6, 4.6, and 3.7 km s 1, respectively. Relative Doppler velocities during the microflare were calculated from the difference between the fitted profile centroids

6 No. 1, 2009 THERMAL AND DYNAMIC EVOLUTION OF A SOLAR MICROFLARE 497 Figure 5. EUV relative Doppler velocities derived from CDS spectra. The average wavelengths against which the relative Doppler velocities were calculated are derived from spectra obtained between 07:30:01 and 07:56:57 UT before the microflare (and before its precursors), from the same 4 20 spatial pixel in which the microflare was observed. We take the uncertainties on the relative Doppler velocities in the He i and O v lines to be the 1σ standard deviations (scatter) in the relevant wavelengths within those time intervals. For He i this yields 1.6 km s 1,andforOv it yields 4.6 km s 1. Horizontal dashed lines indicate ±σ in each frame. The precursors and the microflare itself are delineated with vertical dotted, dashed, and solid lines as in Figure 4. during the microflare and their corresponding rest wavelengths (see Figure 5). Negative values correspond to upward-directed velocities (blueshifts) and positive values to downward velocities (redshifts) RHESSI RHESSI observed the microflare from about 08:11:45 to 08:13:55 UT, but this 2.2 minute event is not included in the RHESSI flare catalog because no emission was detected above 10 kev. (This duration is equal to the lower limit on microflare duration obtained by Hannah et al ) The flare emission peaked at 08:12:16 UT in the higher (7 10 kev) energy range, and it peaked 8 s later (08:12:24 UT) at the lowest (3 5 kev) energies. RHESSI was not observing during the first EUV precursor since its orbit was within the South Atlantic Anomaly (SAA). It emerged from the SAA at 08:03:32 UT, during the second EUV precursor, but observed no emission from that precursor. Figure 6 compares selected EUV, GOES soft X-ray, and RHESSI 3 4 kev and 9 10 kev hard X-ray light curves on an expanded timescale. After the initial EUV burst seen in O v in this figure (and He i in Figure 4, aswell as He ii, Oiii, Oiv, and 1600 Å emission not shown here), the maximum intensities of successively cooler emission occur at progressively later times. See also Table 1. Images of the RHESSI emission show a point source located at ( 584, 142 ), nearly 7 west of the brightest pixel in the 195 Å source. See Figure 2(d). This is the only source detected by RHESSI at the time. 3. RESULTS We describe seven key results that we have obtained for this microflare Precursor Brightenings EUV (Figure 4) and 1600 Å UV light curves reveal two precursor brightenings prior to the microflare itself. See Table 1. After each brightening the intensities return to almost their preprecursor levels. The first precursor lasted 2.4 minutes and appeared only in chromospheric and TR emission, for which intensity enhancement factors relative to pre-precursor averages (07:30:01 07:56:57 UT for CDS, and 07:43:59 07:57:00 UT for TRACE) were 1.33 for He i, 5.20forOiii, 2.35forOv, and 1.29 for 1600 Å emission. The first precursor was accompanied by a hint of blueshifted emission seen only in He i (Figure 5), and it ended about 8.5 minutes before the start of the microflare. The second precursor lasted 6.5 minutes and appeared in chromospheric, TR, and faintly in coronal emission, for which intensity enhancement factors were 1.67 for He i,6.33foroiii, 3.78 for O v, 1.49 for 1600 Å emission, and 1.16 for Si xii. The second precursor ended about 49 s before the start of the microflare. GOES detected no activity during either precursor, consistent with CDS observations that Fe xix emission was absent during these events. RHESSI emerged from the SAA at 08:03:32 UT, during the second EUV precursor, but it too observed no apparent emission from that precursor. Our observations demonstrate that microflares exhibit precursor activity similar to that seen before many larger flares Chromospheric and Transition Region Emission Heralds Start of Microflare The microflare occurred between 08:09:12 and 08:16:03 UT, a 6.9 minutes interval between the systematic rise in the He i and O v light curves (as well as those of He ii, Oiii, and O iv, not shown here) and their subsequent return to local minimum values. See Figure 4. During this period, the cool line intensities

7 498 BROSIUS & HOLMAN Vol. 692 Figure 6. EUV light curves (in ergs cm 2 s 1 sr 1 )ofov at Å, Si xii Å (multiplied by a factor of 8), and Fe xix Å (multiplied by a factor of 20) from CDS, along with hard X-ray light curves (arbitrary units) of 3 4 kev and 9 10 kev photons from RHESSI, plus a soft X-ray light curve (arbitrary units) of Å photons from GOES. The color scheme is indicated in the upper left. Table 1 Timeline of Microflare Events Event Time (UT) Start of 1st EUV Precursor 07:58:16 Peak of 1st EUV Precursor 07:59:44 End of 1st EUV Precursor 08:00:42 Start of 2nd EUV Precursor 08:01:51 RHESSI emerge from SAA 08:03:32 Peak of 2nd EUV Precursor 08:05:56 End of 2nd EUV Precursor 08:08:23 Start of EUV Main Microflare 08:09:12 Max He i ( 16 km s 1 ), O v ( 21) Upflow 08:11:09 Pre-Impulsive Cool EUV Peak 08:11:09 Pre-Impulsive Cool EUV Min 08:11:29 Start Rise He, O, Si Lines 08:11:39 Start Significant RHESSI count rates 08:11:45 Start GOES 1 8 A Emission 08:11:52 Peak of He i, Oiii-v Intensity 08:11:58 Start Rise Fe xix Emission 08:11:58 Peak RHESSI 7 10 kev count rate 08:12:16 Peak RHESSI 3 5 kev count rate 08:12:24 Peak GOES 1 8 A count rate 08:12:30 Peak of Fe xix Intensity 08:12:38 Peak of Si xii Emission 08:13:46 End Significant RHESSI count rates 08:13:55 Last Significant Fe xix 08:14:06 End GOES 1 8 A Emission 08:14:11 End of EUV Main Microflare 08:16:03 Start 1st He, O Post-flare Brightening 08:16:32 End 1st He, O Post-flare Brightening 08:18:59 End Post-flare Transient Brightenings 09:00 began to increase first, then briefly declined again almost as if they were showing another precursor, then began their rapid rise to impulsive maximum. The intensity maximum occurs in only one distinct time element (at 08:11:58 UT, the same for all lines showing the impulsive peak), although the burst extends across two time elements for an apparent duration of about 20 s. Si xii began its rise to maximum (which occurs 1.8 minutes after the impulsive maximum) as soon as the cool emission began its impulsive rise, while Fe xix revealed its first significant increase above the noise at the peak of the cool line intensities. Fe xix emission was observed for a little more than two minutes (see Table 1 and Figure 6), comparable to the lower limit on the duration of microflares derived by Hannah et al. (2008). After the EUV burst, emission maxima in progressively cooler flare plasma (7 10 kev, 3 5 kev, GOES 1 8 Å, CDS Fe xix and Si xii) occurred at successively later times, consistent with cooling thermal flare plasma. See Table 1. Peak enhancement factors (relative to the same pre-precursor average intensities used above) are 3.19 for He i, 24.0 for O iii, 10.7 for O v, and 2.29 for 1600 Å; the (later) maximum enhancement for Si xii was Because the microflare appeared in emission from chromospheric and TR lines before it displayed any increase in coronal or hot flare emission, we conclude that enhanced flare emission at all temperatures is produced predominantly by heating and ionizing the chromosphere; if it were produced at some other location (say, in the corona) then we would observe enhancements in emission formed at that location s temperature before we see it in the chromospheric and TR lines. Initially emission in only the cool lines is enhanced as the chromosphere is heated, but emission at higher temperatures (up to 10 7 K flare emission) is eventually observed as the flare heating mechanism continues to operate. This provides evidence for chromospheric heating by nonthermal electron beams Relative Doppler Blueshifts The microflare s first precursor provides a hint of systematic relative Doppler blueshifts in He i (but not in O v) that correspond to maximum upward velocities of 5 kms 1. See Figure 5, where dashed horizontal lines indicate ±1σ uncertainties on the rest wavelength centroids (1.6 km s 1 for He i and 4.6 km s 1 for O v). The microflare s second precursor was accompanied by more significant, systematic relative Doppler blueshifts with maximum upward velocities of 11 km s 1 in He i and 20 km s 1 in O v. During the microflare itself, He i and O v showed systematic upward velocities with maximum values of 16 and 21 km s 1, respectively, both of which occurred about 50 s (five time elements) before the maximum burst

8 No. 1, 2009 THERMAL AND DYNAMIC EVOLUTION OF A SOLAR MICROFLARE 499 intensity or nearly two minutes after the start of the microflare. Because the Si xii velocity plot does not show systematic trends as clear or well defined as those of He i and O v, wedonot include it here. Nevertheless, we mention that the Si xii line (1σ = 3.7 km s 1 ) yields maximum upward relative Doppler velocities of 8kms 1 around the same time as the maximum upward velocities in He i and O v. In Figures 4 and 6 it can be seen that Fe xix emission is evident for only about two minutes beginning with the maximum burst intensity in the cool EUV lines; Fe xix is not evident during the gradual rise of the cool EUV emission during the microflare s onset, when the cool EUV lines display their maximum upward velocities. Because Fe xix is relatively weak and noisy throughout this interval, we investigated its relative Doppler velocities by generating average profiles during the rise, plateau, and decline phases of its existence. This reveals that the Fe xix emission during its rise phase (08:12:08 08:12:38 UT) was blueshifted about 15 km s 1 relative to its emission during the plateau and decline phases. The velocities observed in He i, O v, and Si xii averaged over this same time interval were 6.8, 6.7, and 2.9 km s 1. Thus none of the other EUV lines showed redshifted emission that would indicate explosive chromospheric evaporation (Fisher et al. 1985a, 1985b; see also Brosius & Phillips 2004; Brosius & Holman 2007) during the period of blueshifted Fe xix emission. Further, none of the He i, Ov, or Si xii lines showed any intervals of significantly redshifted emission during either the precursors or the microflare itself. We conclude that gentle chromospheric evaporation occurred not only during the microflare s precursors, but also during its impulsive rise Electron Density Increase The ratio of the intensity of the O iv line at Å to that at Å is sensitive to electron density (n e ). Its theoretical value varies from at n e = 10 8 cm 3 to at n e = cm 3 (see CHIANTI; Dere et al. 1997; Landi et al. 2006). Based on averaged intensities measured for these two lines in the same time interval used to define rest wavelengths above (07:30:01 07:56:57 UT), we calculate a quiescent electron density of n e = cm 3. (Based on scatter observed in the line intensities during this time interval, we estimate factor of 2 uncertainties on all of the O iv electron densities derived herein.) At the peak intensity in the EUV burst, we obtain n e = cm 3. This maximum density enhancement is short lived, however, as it occurs in only one CDS exposure; values of about cm 3 are found on either side of the peak Compact Source The microflare source seen in Figures 1, 2, and 3 was compact. As described above, the brightening encompassed 11 contiguous pixels in the EIT 195 Å imagery, which corresponds to a solar surface area of about 75 arcsec 2,or km 2.InRHESSI images the microflare appears to be a point source, so its maximum solar surface area corresponds to that of a circle whose diameter is the imager s spatial resolution (3 ), or about 7arcsec 2 ( km 2 ). The microflare emission observed by CDS was confined predominantly to one 4 20 slit pixel. A slight enhancement is detected in the chromospheric and TR emission observed within the next pixel down, consistent with the extension of the TRACE source south and east across that CDS pixel. Magnetograms from MDI reveal that the microflare occurred in and around a negative magnetic island embedded in a larger area of positive magnetic field. This suggests that the microflare occurred in short, low-lying loops that connected the negative magnetic intrusion with its positive magnetic surroundings. Indeed the TRACE 1600 Å contours in Figure 2(c) as well as the TRACE images in Figure 3 appear to be consistent with this idea. Further, if the microflare originated in a unipolar magnetic field region, we might expect that the source observed with CDS corresponds to one footpoint of an extended flare loop. Since neither the extended loop nor its conjugate footpoint are observed in the EIT, TRACE, or RHESSI images, we conclude that the microflare was compact, and confined to small, lowlying loops Evolving Magnetic Field Based on MDI synoptic full-disk magnetograms obtained every 96 minutes, along with a few available full-disk magnetograms obtained at 1 minute cadence, we see that the negative magnetic island evident in Figure 2(b) was growing in size and strength around the time of the microflare. At 06:27 UT, the island comprised two 2 2 MDI pixels with a maximum inward-directed field strength of 96 G. At 08:03 UT (just before the microflare), the island had grown to three MDI pixels with a maximum inward-directed field strength of 100 G. By 08:52 UT it had grown to 12 pixels with a corresponding strength of 210 G, and by 09:35 UT it had become attached to the larger adjacent area of negative field, so that the former island was now more like a peninsula, with a maximum inward-directed magnetic field strength of 283 G. The maximum field strength in the peninsula continued to increase with time, becoming 377 G by 11:15 UT. Thus, the microflare is associated with an area of growing negative magnetic field embedded in a larger area of positive magnetic field No Hard X-Ray Evidence of Nonthermal Electrons Spectral fits to the RHESSI data were limited to the 3 10 kev energy range. The iron line complex at 6.7 kev was apparent, and the count rates were low relative to the background. The spectra were adequately fitted with thermal bremsstrahlung from an isothermal plasma for which the best-fit temperature was found to be highest around the peak of the emission. Using RHESSI s detector 4 alone (to avoid calibration uncertainties experienced by RHESSI at this time), the fit to a spectrum from the interval 08:12:08 to 08:12:28 UT gave a temperature of 13 MK and a volume emission measure of cm 3. These values are close to those obtained from the GOES data, for which the peak temperature is about 12 MK and the corresponding emission measure about cm 3 at 08:12:10 UT. All of these values are well within the range of temperature ( MK) and volume emission measure ( cm 3 ) obtained by Hannah et al. (2008) in their statistical investigation of microflares. Further, using the size of the RHESSI point source to estimate an upper limit on the volume of the thermal X-ray source, we estimate a lower limit of cm 3 on the microflare coronal electron density; this value is greater than the upper limit of cm 3 obtained by Hannah et al. (2008). Because RHESSI spectra could be adequately fitted with thermal bremsstrahlung from an isothermal plasma while they could not be adequately fitted with single or double power-law models, we conclude that RHESSI sees no direct evidence for an electron beam during the microflare.

9 500 BROSIUS & HOLMAN Vol DISCUSSION The EUV light curves and relative Doppler velocities observed during both precursors and the microflare itself are consistent with gentle chromospheric evaporation driven by beamed electrons, presumably accelerated via coronal magnetic reconnection. Because light curves that represent chromospheric and TR emission are the only ones to show intensity increases during the first precursor, no significant coronal heating occurred at that time; if it did, the first precursor would have shown an enhancement in the Si xii or maybe even the Fe xix emission. During the second precursor, intensities of chromospheric and TR lines increased first, before a much smaller increase in the Si xii emission. No significant enhancement in the Fe xix emission was observed during the second precursor. During the actual microflare, intensities of the chromospheric and TR lines systematically increased for 2.5 minutes before the Si xii line began its systematic increase, and 2.8 minutes before the Fe xix emission appeared above the background noise. Further, all of the statistically significant relative Doppler velocities measured in lines formed at all temperatures observed during both precursors and the microflare correspond to upward-directed motion; none of the lines showed significant, systematic redshifts that would indicate explosive evaporation. Although the hard X-ray spectra observed with RHESSI provide no evidence for nonthermal particles during this event, it may simply be that the nonthermal hard X-ray emission associated with the microflare s inferred electron beam is below RHESSI s level of detection. We estimated the nonthermal X-ray emission produced by an electron beam of sufficient energy flux to heat chromospheric plasma to the temperature and emission measure observed by RHESSI, and find that it is below the observed background level for nonthermal energies less than about 10 kev. Thus the RHESSI observations cannot rule out the possibility that nonthermal electrons were produced during the microflare. Milligan (2008) also reported that RHESSI hard X-ray spectra of a GOES B7.6 flare yielded no evidence for nonthermal emission, while Hannah et al. (2008) found that they could obtain a thermal component for more microflares than they could obtain a nonthermal component. We may not be able to rule out the possibility that reconnection occurred lower in the solar atmosphere, possibly even in the chromosphere. If this were the case, then no nonthermal electron beam would have been accelerated in the corona, so no nonthermal hard X-ray emission would have been produced by a beam s stopping in the chromosphere. Indeed, if reconnection occurred in the chromosphere, then the observed gentle chromospheric evaporation would be the result of direct heating of the chromosphere by local reconnection rather than by an electron beam accelerated in the corona. Detailed analyses of additional microflares are required to determine whether or how often direct heating by chromospheric reconnection may occur. Repeated spectral scans of the C iv line at 1548 Å by OSO-8 revealed that redshifts occurred more frequently than blueshifts during C iv bursts (Bruner & Lites 1979; Athay et al. 1980). Since the C iv bursts were sometimes accompanied by a rise of weak soft X-ray emission (Lites & Hansen 1977), some or all of the bursts may have occurred during microflares. The preponderance of redshifted emission observed by OSO-8 during C iv bursts might therefore suggest that explosive evaporation occurs more frequently than gentle evaporation during microflares. However, as Athay et al. (1980) pointed out, in the presence of a rapidly changing intensity (some bursts vary on timescales shorter than the minimum time required to scan the line profile) the emission line s centroid may not be accurately determined because the profile is not instantaneous. It will be interesting to see if a larger sample of instantaneous EUV microflare spectra yield gentle chromospheric evaporation, like we have found in the present work, or explosive evaporation, like that obtained for a FLT reported by Brosius & Holman (2007). One major difference between the microflare reported here and the FLT is that the Sun was quiet during the microflare (with only a GOES C5.9 event at 11:41 UT reported on this date) while it was busy during the FLT, which occurred during (but an arcminute away from) a GOES M1.6 flare. Note that the character of the chromospheric and TR emission changes after the microflare, as evidenced by the numerous transient brightenings similar to precursors that continue until about 09:00 UT. Except for the precursors, the intensities observed before the microflare vary very slowly. Some variation of line intensity in the EUV light curves is attributed to the drift of features (due to solar rotation) across the fixed position of the CDS slit. At disk center the drift rate is about 10 arcsec per hour, so a given point on the disk would take about 24 minutes to drift completely across the 4 slit. It is possible that the brightenings observed after the microflare are due to ongoing occurrences of reconnection (possibly driven by the growing negative magnetic peninsula) that continue to produce spurts of gentle evaporation. Also note a trend toward increasing relative Doppler velocity after the microflare, such that He i returns to near-zero velocities toward the end of the interval displayed in Figure 4, while O v becomes redshifted. This may indicate that material has cooled and is falling back down (as cool rain) after the microflare, but it may be complicated by the persistent precursor-like transient brightenings in the chromospheric and TR emission for nearly an hour after the microflare. This could indicate ongoing reconnection-driven gentle chromospheric evaporation in which sporadic upflowing material attempts to move against the more general downflow. The 7 separation between the EIT and RHESSI sources is greater than the spatial resolution ( 3 ) for either instrument, and therefore appears to be real. The reason for the separation, however, is not certain. Figure 2(a) shows that the RHESSI point source lies between two bright areas of emission observed by EIT, which tempts one to speculate that the RHESSI source corresponds to emission from the top of a loop whose legs are seen with EIT. However, the EIT source to the west of the RHESSI source was observed in both previous and subsequent EIT images, and is not associated with the microflare. It is also interesting to note the location of the RHESSI source in the MDI magnetogram (Figure 2(b)), where it appears in a narrow lane of positive magnetic field that is being pinched by nearby negative fields (especially as the negative magnetic island expands). This suggests an association with reconnecting magnetic fields. It is worth mentioning that there is a time difference (20 30 s) between the EIT image and the hard X-ray peaks around which the RHESSI images were integrated. Further, the RHESSI source temperature (13 MK) is different from that of Fe xii (1.3 MK) or Fe xxiv (20 MK), both of which may contribute to EIT s 195 Å emission during flares. 5. SUMMARY We observed a GOES A5 (after background subtraction) solar microflare over a wide temperature range with RHESSI, TRACE, and three instruments aboard SOHO (CDS,EIT,MDI).Seven

10 No. 1, 2009 THERMAL AND DYNAMIC EVOLUTION OF A SOLAR MICROFLARE 501 key results from this investigation are summarized as follows: (1) the microflare exhibited precursor activity similar to that seen before many larger flares. (2) The microflare brightened in emission from chromospheric and TR lines before it brightened in coronal or hot flare emission, consistent with chromospheric heating by nonthermal electron beams. (3) Measurements of statistically significant relative Doppler blueshifts (with no accompanying redshifts) in EUV lines formed at chromospheric, TR, coronal, and hot flare plasma temperatures indicate that gentle chromospheric evaporation occurred not only during the microflare s precursors, but also during its impulsive rise. (4) A density-sensitive O iv line intensity ratio yields an electron density of cm 3 during quiescent times before the microflare and cm 3 at the peak of the EUV burst. (5) The microflare was compact, covering an area of km 2 in the EIT image at 195 Å, and appearing as a point source in the RHESSI image. (6) The microflare is associated with an area of growing negative magnetic field embedded in a larger area of positive magnetic field. (7) Because RHESSI spectra could be adequately fitted with thermal bremsstrahlung from an isothermal plasma while they could not be adequately fitted with single or double power-law models, RHESSI observed no direct evidence for an electron beam during the microflare. Nevertheless, based on the microflare s observed thermal and dynamic evolution we conclude that it appears to be a miniature flare undergoing gentle chromospheric evaporation, likely driven by beamed electrons accelerated via coronal magnetic reconnection. We cannot, however, rule out the possibility of direct heating by magnetic reconnection in the chromosphere. J.W.B. acknowledges NASA support through SR&T grant NNX07AI09G. 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