I M P U L S I V E P H E N O M E N A IN A S M A L L ACTIVE REGION

G E O R G E L. W I T H B R O E , S H A D I A R. H A B B A L , and R O B E R T R O N A N

Harvard-Smithsonian Center for Astrophysics, Cambridge, 21,1,4 02138, U.S.A.

(Received 24 September, 1984)

Abstract. The temporal and spatial variations of EUV emission from a small growing active region were investigated. Frequent localized short term ( ~ few minutes) fluctuations in EUV emission were observed throughout the 7.2 hr interval when the most continuous observations were acquired. Approximately 20 ~ of the 5" • 5" pixels had intensity variations exceeding a factor of 1.3 for the chromospheric L7 line, a factor of 1.5 for lines formed in the chromospheric-coronal transition region and a factor of 1.4 for the coronal Mgx line. A subflare in the region produced the largest intensity enhancements, ranging from a factor of ~ 2.3 for the chromospheric Lc~ line to ~ 8 for the transition region and coronal lines. The EUV fluctuations in this small active region are similar to those observed in coronal bright points, suggesting that impulsive heating is an important, perhaps dominant form of heating the upper chromospheric and lower coronal plasmas in small magnetic bipolar re~ons. The responsible mechanism most likely involves the rapid release of magnetic energy, possibly associated with the emergence of magnetic flux from lower levels into the chromosphere and corona.

1. Introduction

One of the fundamental problems of solar physics is determination of the coronal heating mechanism. For many years it was widely believed that upward propagating acoustic waves were the primary means by which energy was transferred from the solar convection zone to the corona. However, observations from OSO-8 (Athay and White, 1978a, b) indicated that acoustic waves deliver insufficient energy to the corona to balance the observed radiative and convective energy losses. This finding, coupled with the observation that spatial variations in coronal heating appear to be intimately linked to the strength and configuration of the coronal magnetic field, has led to renewed interest in heating mechanisms based on the dissipation of energy stored in magnetic fields (see reviews by Withbroe and Noyes, 1977; Vaiana and Rosner, 1978; Wentzel, 1981). There are a variety of possible mechanisms for releasing this energy. Some yield quasi-steady heating (e.g. heating by steady currents), while others involve transient, impulsive heating (e.g. rapid reconnection). The objective of the present study was to look for observational signatures of coronal heating processes provided by temporal and spatial variations of the EUV emission from a small active region.

The study was motivated by the results of earlier investigations of the spatial and temporal variations of the UV and EUV emissions from coronal bright points and active regions. Analyses of UV measurements from OSO-8, HRTS, and SMM experiments indicate that UV 'bursts' or impulsive variations in emission from spectral lines formed in the lower chromospheric-coronal transition region (T ~ l0 s K) are a common, nearly continuous phenomenon, especially in active regions (cf. Lites and Hansen, 1977; Bruner and Lites, 1979; Athay et al., 1980; Dere et al., 1981 ; Athay, 1984; Porter et al.,

1984). Similar phenomena have been observed in Skylab EUV measurements of coronal

Solar Physics 95 (1985) 297-310. 0038-0938/85.15 ,i�9 1985 by D. Reidel Publishing Company

298 G . L . WlTHBROE ET AL.

bright points at chromospheric, transition region and coronal levels (Habbal and Withbroe, 1981; Withbroe and Habbal, 1983). The latter observations suggest that the plasma in coronal bright points is heated by an intermittent, impulsive heating mecha- nism, most likely involving the rapid release of magnetic energy. The results discussed below suggest that impulsive heating processes are also important in small active regions.

2. The Observations

The EUV observations used in the present investigation were acquired by the Harvard Skylab experiment (see Reeves et aL, 1978a, b). The data consist of 5' • 5' spectro- heliograms made in HI L~ 2 1216, CII ). 1335, Cm 2977, OIv 2554, OvI 2 1032, and Mgx 2 625. Measurements were made simultaneously at these six wavelengths which span the temperature range from the chromosphere (~ 2 x 104 for L~ and C n) to the c o r o n a (10 6 to 2 x 10 6 K for Mgx). The three transition region lines are formed at mean temperatures of approximately 105, 2 • 105, and 3 • 105 K, respectively for Cm, Ow, and O vI. The spatial resolution was 5" and temporal resolution 5.5 min. The instrument employed a photoelectric detection system.

Figure 1 contains a set of spectroheliograms in the above lines made near Sun-center. East is toward the top, north to the right. One of the coronal bright points (No. 6) studied by Habbal and Withbroe is circled. The dark areas surrounding this bright point and extending downward to the left in the Mgx spectroheliogram are part of a large coronal hole that extended across center of the solar disk at the time of the observations. The active region investigated in the present study is the bright patch of emission (see Mgx image) immediately below the circled bright point. In the lower temperature lines the active region is split into two areas with the eastern (upper) area somewhat larger and brighter than the western (lower) area at the time these images were acquired.

3. Temporal Variations

Figure 2 illustrates the temporal variations in the integrated emission from the two halves of the active region. The left-hand plots are for the eastern patch of emission, which is the largest and brightest area of emission in Figure 1, while the right-hand plots are for the emissions from the western part of the active region. The plots are arranged in descending temperature of formation from Mg• in the corona to L~ in the chromo- sphere. All of the fluctuations are statistically significant. (The uncertainties are smaller than the plotted points.) One unit on the logarithmic intensity scale corresponds to 10 0.04 o r about 10% on a linear scale. The largest peak shortly after 06 : 00 UT in the eastern half of the region is due to a subflare. Based on differential emission measures estimated from the data (see Habbal and Withbroe, 1981), the radiative energy released by this subflare was 1028 erg. Smaller fluctuations with energies of the order of 10 26 erg occurred nearly continuously in both halves of the active region.

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Fig. 2. Left-hand plots: Variation of the total intensity in the eastern area of the small active region as a function of time for L~ (chromosphere), C I[I, O Iv, and O vI lines (transition region) and Mg x line (corona). The expected variations due to statistical fluctuations are smaller than the size of the data points. The dashed lines connect points during long intervals when there were no observations (usually due to spacecraft night). The intensity peaks at 05 : 24 and 06 : 03 UT are due to subflares visible in Figure 5. One unit on the ordinate corresponds to 0.04 dex ( = 1 0 ~ 1 7 6 Right-hand plots: Total intensities from the western area of the

active region.

I M P U L S I V E P H E N O M E N A IN A S M A L L A C T I V E R E G I O N 301

Note in many cases that there is a strong correlation between brightenings observed at chromospheric through coronal levels. Since the M g x is a fairly 'cool' coronal ion with a maximum fractional concentration at logT.~ 6.15, one might expect that Mg x 2 625 could be sensitive to heating or density fluctuations occurring in the tran- sition region. We investigated this and found less than 4~o of the Mg x emission originates from material at temperatures below 106K for plasmas with differential emission measures similar to that of the observed region. Consequently, given the magnitude of the fluctuations observed in Mgx and in the transition region lines, the Mgx fluctuations must occur in plasma with temperatures greater than 10 6 K. Thus, the data in Figure 2 indicate that there is a significant 'ac' component to the emissions from the active region originating from plasmas at chromospheric to coronal tem- peratures. Since these curves are for the integrated emission from the two halves of the active region, the curves are averages of larger amplitude (AI/I where I is the intensity) fluctuations in the fine scale structure of the region.

In order to obtain information concerning the range and frequency of the observed fine scale intensity fluctuations, we determined for each spectral line the fraction N/Ntot~ 1 of the 5" • 5" pixels in the active region whose intensity changed by a factor of F or more in 5.5 min. The resulting variation of N/Ntota ~ with F is plotted in Figure 3. The subset of the data used for this aspect of the study consists of measurements of intensities of 266 pixels in the active region. The intensity of each pixel was measured 35 times over a period of 7.2 hr. Figure 3 shows that the largest enhancements are detected in the transition region lines and the smallest in the chromospheric L~ line. Approximately 20"/o of the pixels had intensity variations exceeding a factor of 1.3 for L~, a factor of 1.5 for lines formed in the transition region and a factor of 1.4 for the coronal Mg x.

Some of the observed intensity fluctuations are caused by statistical fluctuations in the measured intensities. In order to minimize the effects of statistical fluctuations we used only those pixels whose intensities were greater than a selected threshold whose value depended on the mean intensity (in counts) of the spectral line. The adopted thresholds were 200, 40, 200, 80, 80, and 20 counts per integration period respectively for the L~, CII, CIII, OIv, OvI, and Mgx lines. This eliminated approximately 35o./0 of the pixels for C II, O IV, and Mg x and 10 ~/o of the pixels for the other three lines. Given the magnitude of the adopted intensity thresholds, one can calculate for each line the maximum contribution of the statistical fluctuations to N/Ntota ~. This contribution falls off much more rapidly with F than the observed values of N/Ntota~. For example, for the weakest line, Mg x 2 625, the calculated contribution of statistical fluctuations falls to a value N/Ntota ~ smaller than 0.07 by F = 1.4 and smaller than 0.006 by F = 1.6. For the much stronger L~ line the calculated contribution of statistical fluctuations falls to less than 0.005 by F = 1.2.

In order to provide insight into the amount of energy contained in the fine-scale short-term intensity fluctuations, we determined for each line the amount of power in intensity changes of amplitude AI in 5.5 rain, the time interval between successive spectroheliograms. Measurements from all pixels in the region were summed over the

302 G. L. WITHBROE ET AL.

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7.2 hr interval when the best observations were acquired. From these data one could then estimate the energy radiated in the fluctuating component of the emission from the region. For L~ the radiated energy was calculated directly from the observed L~

I M P U L S I V E P H E N O M E N A IN A S M A L L A C T I V E R E G I O N 303

intensities. For calculating the energy radiated by the transition region and corona we

utilized differential emission measures estimated from the data (see Withbroe, 1977,

Habbal and Withbroe, 1981) and the radiative loss function given by Rosner et al.

(1978b). The results are plotted in Figure 4. The curves give the fraction f o f t h e energy

changes greater than the value AE,. (in units of E,., where Er is the mean energy radiated by the active region, 1.2 x 1025 erg s - l in L7 and 1.0 x 1025 erg s i in the transition

region and corona). F rom the curves in Figure 4 we can calculate the fraction of the energy radiated by

the active region which is variable on a time scale of 5.5 min. For the chromospheric

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transition region and corona.

304 G. L. W I T H B R O E ET AL.

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Fig. 5. Contour maps constructed from M g x 2625 spectroheliograms showing some of the spatial and temporal changes in the small active region. The time (hr, min) that each spectroheliogram was acquired on August 21, 1973 are given. East is to the left, north to the top. The contours were drawn at levels of 60, 90, 150 (shaded light gray), 230 (shaded dark gray) and 350 (shaded black) counts/integration time. The mean background level was approximately 2 counts in the coronal hole areas and 30 counts in the quiet regions (areas not in bright points, active regions or coronal hole). The contoured area is 100" x 55" with the tick marks at intervals of 5" (horizontally) and 2.5" (vertically). This area is in the upper right quadrant of Figure 1 (note rotation of 90 degrees, since east is at the top of Figure 1) immediately below bright

point 6.

I M P U L S I V E P H E N O M E N A IN A S M A L L ACTIVE R E G I O N 3 0 5

L~ line this fraction is ~ 15 O//o and for the energy radiated by the transition region and corona this fraction is ~ 30 %. These values are lower limits to the magnitude of the short term ( < few min) 'ac' component of the energy radiated by the region due to the limits imposed by the spatial resolution (5") and temporal resolution (5.5 min) of the obser- vations. Hence, it appears that a substantial fraction, perhaps all, of the energy radiated by the fine scale structure of the active region varies on time scales of a few minutes. This implies that a substantial fraction of the energy input to these layers has a short term 'ac' or impulsive component.

4. Spatial Variations

The nine contour maps in Figure 5 show the spatial development of the Mg x emission at intervals of 5.5 min from 05:24 to 06:09 UT. The intensity difference between adjacent contours is a factor of 1.5. (These maps are rotated 90 degrees compared to Figure 1 and have east on the left.) The large enhancements in intensity in the top and bottom images are caused by subflares. Note that the smaller subflare at 05 : 24, which is very prominent in the contour maps, produces only a small peak in the integrated Mg x intensity in Figure 2. Also note the changes in the emission in the lower left-hand comer of the active region. Some of these changes appear to be associated with the one or both subflares. During the time period covered by Figure 5 the right-hand (western) area of the active region had weak emission and only small brightenings (note the increase and then decrease in the area of the second contour level of intensity between 05 : 46 and 06 : 09 UT). Also, the coronal emission connecting the eastern and western areas was too weak to show up at the first contour level.

Figure 6 shows lower level fluctuations (compared to subflares) in the intensity of the eastern (left-hand) area of the active region. Note that the locations of intensity changes vary from image to image. Also note the development of the western (right-hand) area of the active region which increases in area and brightness from the first image to the last (as illustrated by the increased area of the second level contour). By the time of the next orbit (Figure 7) the western or right-hand area has increased significantly in intensity (3rd contour level) followed by a slow decline. A bridge of bright Mg x emission connecting the eastern and western areas became prominent. (This bridge of emission appears to be from coronal loops connecting the two regions, since the eastern and western areas are not connected by bright emission when observed in spectral lines formed at chromospheric and transition region temperatures - as shown in Figure 1.) Note the fluctuations in brightness in various parts of the active region (e.g. appearance and disappearance of 2nd contour at lower left in maps for 08 : 27-08 : 38; appearance at 08 : 55 and disappearance of small 3rd level contour in the central part of the eastern (left-hand) area).

Figure 8 shows the longer term coronal development of the active region over a period of approximately 12 hr. Because the pointing of the instrument was fixed, the active region moved across the field of view from the eastern (upper edge) of the spectro- heliogram at 22 : 21 UT on August 20 to approximately the middle of the last spectro-

306 G. L WITHBROE ET AI..

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heliogram at 10:46 on the next day. Note that the active region grows in area and brightness over most of this time. The western part of the region does not appear with significant intensity until third frame (03 : 01 UT). Also note how the coronal emission

bridging the eastern and western areas of the active region becomes more intense as time progresses.

IMPULSIVE PHENOMENA IN A SMALL ACTIVE REGION 307

. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

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Unfortunately the 5" spatial resolution of the observations is insufficient to resolve individual loops. We cannot determine from the observations whether magnetic recon- nection played a role in the observed structural changes (or to be more specific, the observed spatial variations in EUV brightness). Given the rapid growth of the region during the period of observations (see Figure 8), it appears likely that emergence of

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I M P U L S I V E P H E N O M E N A IN A S M A L L A C T I V E R E G I O N 309

magnetic flux from below the surface played a significant role in these brightness/struc- tural changes. It is interesting to note several coincidences: (1) that the western section of the active region began growing after the subflare activity and subsequent decline in transition region brightness in the eastern section (see Figure 2). (2)Similarly the enhancement of the brightness of the loops connecting the eastern and western sections of the active region occurred after the large jump in EUV emissions from the western section and during the subsequent decline of these emissions (see Figures 2, 5-8). One might speculate that as magnetic flux elements emerge from the surface they are stressed and that the release of these stresses produces heating which occurs during or shortly after emergence.

5. Summary and Conclusions

The EUV emissions from the small active region investigated in the present study exhibit significant variations over temporal intervals of 5 to 15 min in structures with spatial scales of 5" to 15". At the limit of the temporal resolution, 5.5 min, there appears to be nearly continuous fluctuations in the emission from lines formed at upper chromo- spheric, transition region and coronal levels. Longer term evolutionary trends are also observed, such as the growth of the eastern and then western halves of the region and the enhancement of loops connecting the two halves.

These observations of a small active region and the results of analyses of EUV observations of coronal bright points (Habbal and Withbroe, 1981; Withbroe and Habbal, 1983), suggest that stochastic processes play an important, perhaps dominant role in the heating of the upper chromospheric, transition region and coronal plasmas in small scale (< 1') magnetic bipolar regions. These heating processes most likely involve the rapid release of magnetic energy, possibly associated with the emergence of magnetic flux. Energy could be stored in the magnetic field by twisting of the field lines into non-potential configurations before emergence from the photosphere and/or by systematic twisting produced by turbulent motions in the photosphere after the field lines have emerged (cf. Tucker, 1973; Rosner et al., 1978a; Parker, 198 la, b, 1983a, b). The following hypothesis is suggested. It is a scenario similar to that which has been suggested for flares in active regions (see review by Sturrock, 1980). The stored magnetic energy is released suddenly (via for example rapid magnetic reconnection) resulting in plasma heating, most likely at coronal levels. Energy is transferred to the chromospheric footpoints of the magnetic loops via thermal conduction and/or accelerated particles. This deposition of energy into the chromosphere heats up the upper chromospheric material which flows upward filling the loop with hot, dense plasma via 'chromospheric evaporation'. There is a 'burst' of EUV emission lasting several minutes due to the impulsive plasma heating. After the impulsive heating phase, some, perhaps most of the plasma drains from the loop, depending on the amount of 'steady state' heating, if any, continuing after the impulsive phase. The Doppler shifts measured in impulsive brightenings observed in UV lines formed at T -~ 105 (see Athay (1984) and references cited therein) could be produced by the upward and downward flows of heating and

310 G.L. WITHBROE ET AL.

cooling material in the loops, with the relative proportion of observed upflows and downflows depending on the relative magnitude of the differential emission measures of the heating and cooling plasmas. Waves associated with the heating process could also produce some of the observed Doppler shifts.

Measurements with higher spatial and temporal resolution are required in order to determine whether the 'dc' component of the EUV emission is the result of heating by a slowly varying steady-state process or an integration over many small impulsive events. At the 5" spatial resolution and 5.5 min temporal resolution of the EUV Skylab observations we may be seeing only the larger impulsive events and missing those occurring on smaller spatial and temporal scales. Observations with more sensitive instruments having better spatial and temporal resolutions are required. Of particular interest are instruments to be flown on Spacelab 2 which are capable of acquiring high resolution magnetic field measurements in the photosphere and high resolution measurements of spectral line intensities and profiles of lines formed in the upper chromosphere and lower transition region.

Acknowledgements

We thank D. Wentzel for his helpful comments and suggestions. This work was supported in part by NASA under grant NAGW-249 and the USAF under contract F19628-82-K-0018.

References

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