solar flares in the extreme ultraviolet

S O L A R F L A R E S I N T H E E X T R E M E U L T R A V I O L E T

L, A. H A L L

Air Force Cambridge Research Laboratories, Bedford, Mass. U.S.A.

(Received 25 June, 1971)

Abstract. Measurements of flare-related impulsive enhancements in solar emission lines in the extreme ultraviolet, observed from the satellite OSO-III, are reported. The enhancement of a line, expressed in percent of the total disk intensity in the line, is of the same order of magnitude as the flare area, expressed in heliocentric square degrees. Rise-times and decay-times of impulsive enhancements average about 2 min and 5 min, respectively. The maximum enhancements of radiation from ions in the chromosphere-corona transition region precede the Ha maximum by an average of 2 min, and occur in the same period of time as the hard component of solar X-rays and the impulsive microwave bursts. Coronal lines in the extreme ultraviolet are less impulsive than the transition region lines in flare-related enhancements and their maxima follow the Ha maximum.

1. Introduction

With the advent of rocket and satellite technology, it became possible to observe the

sun in a hitherto inaccessible spectral region, the extreme ultraviolet (EUV). One such

set o f observations was made by the A F C R L spectrometer aboard the N A S A satellite

OSO-III . A tabulat ion o f the solar flux in the range 1310-270 A and the variat ion

with solar activity of a number o f emission lines in this range has been published (Hall

and Hinteregger, 1970). The instrument was also operated extensively in a mode well

suited to the observation o f solar flares in a single emission line with high time reso-

lution. A few of these observations have been reported (Hall and Hinteregger, 1969)

and the subject of this paper is a more complete survey.

2. Measurements

The AFCRL instrument aboard OSO-III was a grating spectrometer measuring the solar spectrum in the range 1310-270/k. A description of the instrument has been given elsewhere (Hall and Hinteregger, 1970). The pointed section of the spacecraft kept the optic axis of the spectrometer trained on the center of the disk continuously throughout the sunlit portion of the orbit. The entrance aperture of the spectrometer was designed to admit light from all parts of the solar disk without discrimination, so that at all times the measurement was of the total integrated light from the sun. The spectral resolution of the instrument was 1.5/k at the short wavelength end of the spectrum, increasing monotonically to 3.5 A at the long wavelength end.

Two major modes of operation were provided, a scan of the total wavelength range requiring 5.44 min for completion, and a fixed-wavelength mode wherein the instrument scan had been stopped by ground command at one of the 2040 steps of the

Solar Physics 21 (1971) 167-175. All Rights Reserved Copyright �9 1971 by D. Reidel Publishing Company, Dordrecht-Holland

168 L.A. HALL

total scan range. Since the instrument had one detector, only one line was observed at a time in either mode of operation. In the scan mode, therefore, a given line was observed every 5.44 min, except for cases where the line of interest appeared in more than one order of diffraction. In the fixed-wavelength mode, measurements of the wavelength under observation were repeated at intervals of 0.16 s, the read-out rate of the instrument into the tape recorder. In this mode, therefore, the variation of intensity of a solar line could be observed in high time resolution. This mode was well suited to two different types of measurement, the absorption analysis of the terrestrial atmosphere by monochromatic probing at satellite sunrise and sunset (Hinteregger and Hall, 1969), and the observation of solar flares in the extreme ultraviolet emission lines, which will be described below.

Since the instrument gave a measurement of the integrated light from the whole disk, the variation caused by a flare in a small region of the disk was of the order of a few percent, even if the intensity in the flaring region increased by orders of magnitude. The statistical accuracy of the observations in the stronger lines, however, was such that an enhancement of signal of only one percent could be detected. The statistical accuracy of the measurements declined over the life of the experiment as the sensitivity of the instrument irreversibly decreased, presumably because of progressive gain loss in the magnetic electron multiplier used for a detector. The measurements reported here span the period 10 March 1967 to 1 August 1967, after which the instrument performance discouraged further search. In that period observations in the fixed-wavelength mode have been examined for enhancements during a total of 340 flares of all sizes. In about one-third of these cases enhancements in the EUV were observed. Most of the flares for which no enhancement could be detected were sub- flares and, as will be seen, the relation of EUV enhancement to the area of the flare showed that, in most wavelengths, enhancements during sub-flares would be expected to be below the level of detection in the instrument used.

Apart from the progressive sensitivity loss, there were other limitations. A major difficulty is that every solar flare is different and each of them was observed in only one of the 16 different wavelengths monitored in the fixed-wavelength mode. Therefore, in what follows, the conclusions are necessarily arrived at on statistical grounds; that is, some sort of average or typical behaviour is inferred from a sample in which the dispersion of the measurements is uncomfortably large.

Another limitation is an observational bias in favor of the impulsive event. Experi- mental difficulties prevented reliable detection of a long, slow increase in the solar flux; for example, an increase of a few percent over a time span of 30 min or more.

Figure 1 gives examples of observation in four different wavelengths, showing some of the variety encountered. The event depicted in Figure l(b), in the 304 N line of Hen, is unusually beautiful because it is a very large event occurring in an intense line so that the statistical accuracy is good. All except the smallest fluctuations in the curve are real variations in intensity. The general shape of this enhancement is seen repeatedly in the different lines, singly or in superimposed forms with more than one peak. Even the initial slower rise before the main fast rise is observed in many cases.

SOLAR FLARES IN THE EXTREME ULTRAVIOLET 169

2 4 0 0

2 0 0 0

(o)

20 MARCH 1967 1206.5A" (S[ "nT)

l 08150 I 09 IO0 I 09110 q 0920l I .

UNIVERSAL TIME

~ I 0 , 0 0 0

5000

0 4 5 0

(b)

2 6 M A R C H 1 9 6 7

303.8 A (He 'f'r)

u I ~ / t I 0 5 0 0 05 fO 0 5 2 0

UNIVERSAL TIME

4000

5500!

5 0 0 0

2 5 0 0

(c)

3 0 M A R C H 1 9 6 7

6 2 9 . 7 A (Ox]" )

0 0 0 0 0 0 1 0 0 0 2 0 0 0 3 0 UNIVERSAL TIME

(d)

t A P R t , L S' S', '

1 0 3 1 . 9 A ( 0 " ~ 1 1,

2 0 0 0 o

< li m 1500

0610 0 6 2 0 0 6 3 0 0 6 4 0 0 6 5 0 UNIVERSAL TIME

Fig. 1. Some flare-related EUV enhancements observed from OSO-III. In plots (a) and (c), the drop at later times is due to eclipse by the Earth. In plot (d), there is a superimposed oscillation of about 3.3 min period and 2 % magnitude which is instrumental, not solar.iThe data have been~averaged in

blocks of 30 samples (4.8 s).

3. Results

A. RELATION OF EUV ENHANCEMENT TO FLARE AREAS

In order to arrive at some comparison between EUV observations in a variety of

wavelengths made on different flares, it is necessary to rate flares in some way. The

most widespread custom is to assign ' importance ' by means of flare area observed in

Ha, even though the dispersion in a set of measurements of the same flare by different

observatories is considerable. For example, in the set of 27 flares observed in the

630 A line by our OSO-III spectrometer between 20 March 1967 and 5 August 1967,

the scatter in the reported values of the area of a particular flare (Solar-Geophysical Data, 1967) is typically about +_40%. Presumably the average adopted by the com- pilers is somewhat better than the individual measurements, but the number of

observatories reporting is never large enough to gain much additional accuracy by averaging the reported values. For the present purpose the areas used were the average measured areas reported in Solar-Geophysical Data, corrected for position on the

solar disk by using the relation (Sawyer, 1967):

a A =

cos 0 + 0.2 sin 0

170 L.A.HALL

where A is the corrected area, a is the measured apparent area, 0 is the heliocentric

angle from the center of the disk. For central distances less than 0.4 (0 < 24~ the

correction reaches only 2~, which is trivial compared to the errors already present,

and the correction was therefore not made.

Almost one-third of the fixed-wavelength measurements were made in the 304 A

and 630/~ lines, because these lines were especially useful in the absorption measure-

ments at satellite sunrise and sunset. Preliminary inspection showed that enhancements

of about the same magnitude were observed in these lines for the same flare class.

Therefore these measurements were lumped together for the purpose of attempting to

determine the functional relationship of their enhancements to the flare areas. For

the 55 events in which flare enhancements were observed in these lines, a plot of

percentage enhancement, E, versus flare area, A, showed that the data points fitted

best to a line E ~ A 3/a. The fit to this line, however, was only slightly better than the

fit to a line E ~ A , because of the scatter in the data points. This scatter is of the same

order as the errors introduced by the uncertainty in the area measured in He, and this

may be the dominant error.

For the sake of uniformity in handling the data the A 3/2 dependence was assumed

for all the lines. It cannot be assumed however that a 'best-fit' to a line E ~ A 3/2 im-

plies that the enhancement is truly proportional to the volume and that the flare is

therefore optically thin in the radiation. This may be true, at least for some lines, but

there are many other ways in which an apparent 3/2-law can be explained provided a

suitable model is invoked. For instance, the intrinsic brightness in the EUV may be

larger in the larger flares, which would cause radiation from even an optically thick

source to increase as A", with n greater than unity.

Table I gives in column four the average values of the constant of proportionality,

TABLE I

Relation of impulsive EUV enhancements to flare areas

logzo Te Ion Wavelength k No. of events (angstroms) (average) entering average

4.0 H 1215.7 0.3 ~ 0.1 2 4.0 H 1025.7 0.5 • 0.2 3 4.0 H 972.5 0.7 • 0.2 4 4.0 H cont. 880 1.7 • 0.6 5 4.3 He~ 584.3 < 0.3 3 4.5 Sire 1206.5 2.4 4- 1.6 9 4.7 HelI 303.8 1.4 • 0.8 23 4.7 Cm 977.0 0.9 E= 0.5 2 5.3 Ov 629.7 1.2 ~ 0.4 26 5.5 OvI 1031.9 2.8 :t_ 1.4 13 5.8 Nevm 770.4 1.0 =E 0.4 4 6.0 MgIx 368.1 < 0.5 5 6.1 Mgx 625.3 < 0.3 5 6.3 SixlI 499.3 < 0.5 8 6.3 Fe xv 284.1 < 0.4 3 6.4 FexvI 335.4 < 0.4 2


k, in the equation E=kA 3/z, for each wavelength listed. Column five gives the number

of observations averaged together to arrive at the value of k in column four. The error

assigned to the k-values is a simple mean-deviation-from-the-mean, since the number of events is not large enough for any more sophisticated analysis. The k-values pre- ceded by the 'less-than' symbol are cases in which no impulsive enhancements were observed and the upper bound is arrived at by an estimation of the limit of observa- bility in the measurement process.

Column one of Table I gives the electron temperature at which the emitting ion species is favored in the undisturbed solar atmosphere (Jordan, 1969; Cox and Tucker, 1969), and serves primarily to order the data in a relative way. The conditions of formation and excitation of the emitting ions in a flare are of course much different from those assumed in the equilibrium calculations leading to the temperatures listed.

The most obvious feature of the data in Table I is that the values of k are all of the same order of magnitude for the ions formed at temperatures less than 106 K. (An apparent exception is the 584 • line, to be discussed later.) The scatter in the individual measurements entering into each_ average, as represented by the mean deviations given, is of nearly the same magnitude as the apparent differences between the average values for the different ions, which discourages speculation as to what the differences may mean. However, the lower values of k for the hydrogen lines, and the trend of their values, leads one at least to suspect that differences in optical thickness among these lines may be involved.

To put the results of Table I in terms of the more familiar importance classification of flares, a k-value of unity for a line means that on the average its intensity integrated over the total disk will increase about 1~o for a sub-flare, 7~ for a flare of Class 1, and 25~ for a flare of Class 2.

The upper bounds listed for the ions formed at temperatures greater than 106 K are consistent enough and based on enough observations to suggest that they delineate a real difference between those ions and the ones formed at lower temperatures. In this connection it must be repeated that the data in Table I represent the occurrence of impulsive events, seen in the fixed wavelength mode of the instrument. Some flares were observed in the scanning mode, and there is evidence from these scattered observations that the last five lines in Table I sometimes at least have a gradual rise and fall in intensity late in the flare, on the order of ten minutes after the impulsive events (Hall and Hinteregger, 1969).

In some of these flares occurring while the instrument was scanning, enhancements were observed in the 584 Zt line, similar in magnitude to those of the nearby 630 A line, so that the value of k entered in Table I for the 584 ~ line (based on only three events) is probably not typical. In these observations in the scanning mode, the time resolution is 5.44 min, so that there is no good way to determine the magnitude of the peak enhancements and thereby arrive at k-values for Table I.

The observations in the scanning mode, however, do show that there are lines in the solar spectrum exhibiting unusually large enhancements. These are the triplet lines from the Be-like ions CIII and Ov and the Mg-like ion SiliI. They show enhancements

172 L.A. HALL

about three times as large as those of the singlet resonance lines of the same ions (Hall and Hinteregger, 1969). Unfortunately, none of these triplet lines were chosen

for observation in the fixed-wavelength mode, and we therefore have only the low time resolution measurements. These triplet lines are emitted following collisional

excitation from a low-lying metastable state, the population of which is also dependent

on collisional excitations. The step-wise nature of the process populating the upper

triplet state leads to an emission intensity which is more sensitive to electron density

than is the singlet resonance line. The singlet-triplet ratios of these ions have been observed to change with solar plage activity in non-flare conditions (Hall and Hinter-

egger, 1970; Munro et al., 1971) and the ratios may be used to derive electron densities (Munro et al., 1971).

B. TIME-DEPENDENCE OF ENHANCEMENT

The term 'impulsive' in the preceding parts of the discussion can be delimited more

precisely by considering the average time constants of the observed enhancements. Previous reports on the data from the A F C R L experiment on OSO-III (Hall and

Hinteregger, 1969) and the Goddard Space Flight Center experiment on the same

satellite (Neupert, 1969) have established that the EUV maximum commonly occurs

some minutes before the Hc~ maximum. Even earlier, Donnelly (1967) had concluded

from the analysis of a type of ionospheric disturbance known as a 'sudden frequency

deviation' (SFD) that it was caused mainly by EUV radiation, and if so, the EUV flash occurred prior to the Hc~ maximum. The present study attempts to define the

time dependence more precisely by a statistical analysis of about 50 well-timed EUV enhancements in various wavelengths, observed f rom OSO-III.

Columns four and five of Table I I give the average values of the rise-time and decay-

time of EUV enhancements for several ion species. (In Table II, the hydrogen line and

continuum observations have been averaged together to improve the accuracy of

their average time constants.) The rise-times and decay-times in the table are not time

constants in the more usual sense of an e-fold rise or decay, but are the total times

measured from the first detectable rise from the pre-flare level to the maximum, and

TABLE II Time constants of impulsive EUV enhancement

log10 Te Ion Wavelength Rise-time Decay-time At(Hc~ max. (angstroms) (rain) (rain) --EUV max.) (rain)

4.0 H Lyman series 2.4 -4- 1.6 (6) 4.5 Silii 1206.5 2.3 4- 1.7 (7) 4.7 HeII 303.8 2.4 4- 1.3 (14) 4.7 Cm 977.0 1.1 • (4) 5.3 Ov 629.7 1.5 d= 1.0 (19) 5.5 Ovr 1031.9 2.0 • 1.1 (15) 5.8 Nevm 770.4 1.5 • 0.5 (2) 6.4 FexvI 335.4 8 (1) < 6.0 Averages 2.0 • 1.2 (67)

4.4 4- 3.0 (5) 3.4 • 0.9 (4) 4.7• (7) 4.2• (5) 5.9• (9) 2.1.4.0.9 (9) 2.3-4-0.2 (3) 1.4.4.0.9 (3) 4.7-4-2.5 (16) 2.1,4,1.6 (13) 5.7 ~: 3.6 (12) 2.0 -4- 1.4 (5) 7.2 • 2.8 (2) 0.5 (1) 19 (1) -- 8 (1) 5.0 • 2.7 (54) 2.4 • 1.5 (40)


from maximum back to the initial level, respectively. Some of the decay times were

hard to establish, since the decay is more gradual than the rise; therefore the rise-time

averages in the table are based in most cases on more events than the decay times. The

number in parentheses following the average is the number of observations entering into the computation of that average. Again in Table II, the error limits given are in the form of a mean-deviation-from-the-mean, and indicate the extent of the variation

in time constants from one event to another. This dispersion in the values entering into the average prevents a definite conclusion as to whether or not there is a real trend toward higher or lower values as one looks through the temperature range represented. There is an apparent trend toward shorter rise-time and longer decay-time as one goes to the higher temperatures of formation. Simple averages over all events for the lines from ions formed at temperatures less than 106 K are given in the last line of the table, i.e. rise-time 2.0 __+ 1.2 min and decay-time 5.0 + 2.7 min.

The sixth column in the table gives the time interval by which the EUV maximum in a particular line precedes the He maximum as reported in Solar-Geophysical Data. The dominant feature of the data in column six is that the maximum EUV enhancements precede the He maximum by an average of 2.4 min for those lines normally emitted from regions of temperature less than 106 K. The dispersion in the values entering the average for a particular line is represented in the table by the mean-deviation- from-the-mean. The uncertainties are unfortunately just too large to permit a firm conclusion, but the data seem to show a later EUV maximum as one goes to higher temperatures. In fact, one ma~y discern a grouping in the data in column six as follows. The lines of hydrogen and Sire precede the He maximum by three to four minutes. The lines of Hell, Cm, Ov, and OvI precede the He maximum by about two minutes. The lines of Ne vm and Fe xvI have enhancements following the He maximum, and the same behaviour has been observed in lines of Mgx, Sixn, and Fexv in flares observed while the instrument was in the scanning mode. The middle group includes ions over a wide temperature range, but as some current solar models show (Athay, 1966; Dupree and Goldberg, 1967), these ions are formed in the narrow transition zone around 2000 km above the limb where the temperature rises from 4 x 104K to 2.5 x 105 K over a height range of only about 200 km. It is therefore not surprising that their enhancements are virtually simultaneous. The ions in the first group, log T~<4.5, are formed 1000 km or more lower in the model, and the ions in the third group, log T>~ 5.8, are coronal, formed at altitudes of around 3000 km and up. If the different excitation times are simply the result of an outward motion of the exciting agent, rather than intrinsic differences in rate-coefficients for the processes in the different species, the velocity of the excitation agent in the model cited would be of the order of 10 km/s. At least one immediate objection to the foregoing speculation (besides the limited accuracy of the observations) is that the model used is derived from observations of the quiet sun, not the region above active centers, and is further- more a model of a stratified atmosphere. However, the observation of near-simul- taneity of the enhancement of the ions of the middle group lends support to the con- cept of a narrow transition region, whether in a stratum or in spicular material.

174 L.A. HALl.

The last line entered in Table II represents a single case of an enhancement in the 335 A line of FexvL a clearly coronal ion. As the much longer time-constants show, it does not fit into the category heretofore called impulsive, and it occurred well after the He maximum. This is the single clear observation of an enhancement of this line in the fixed wavelength mode of the instrument, out of 12 flares where an enhancement might have been observed in this line. However, corroborative observations were made in the large flare around 0022 UT of 22 March 1967, while the instrument was scanning in wavelength. This 335 A line of FexvI and emission lines of NevnI, Mgx, SixH, and Fexv were found to be enhanced about 10 or 15 min after the He maximum (Hall and Hinteregger, 1969).

4. Correlation with Other Observations

In the foregoing discussion, the EUV impulsive enhancements have been related in their magnitude and timing with He flare observations, simply because the He patrol provides the most extensive and easily available set of observations. The consistent appearance of the EUV maximum prior to the He maximum by 2 min or so places the EUV maximum in the flash phase of the flare, along with hard X-rays and impulsive microwave bursts (Kane, 1969; Kundu, 1965). A survey by Castelli and Richards (1971) of those of our EUV enhancements occurring during the AFCRL patrol of the solar radio noise in centimeter wavelengths shows that of 30 associated events, 70~ showed EUV and radio maxima occurring within one minute of each other. In the meter wavelengths, the Type III burst shows a high degree of correlation. The Type IIl events listed in Solar-Geophysical Data were examined and of 41 EUV events with enhancements greater than 2%, occurring in periods covered by the radio patrol, 20 had associated Type II1 bursts. For 18 of these associated events, the complete time span of the burst was reported, and in 15 cases the EUV maximum fell within the time span reported, even for 5 Type III bursts of 1 or 2 rain duration.

Early in the examination of our flare-related enhancements it became obvious thai a close relationship exists between the EUV enhancements and sudden frequenc5 deviations in radio carrier frequencies reflected in the E and Fregions of the ionosphere (Donnelly, 1969; Hinteregger and Hall, 1969). On the basis of this relationship, Donnelly has used his observations of sudden frequency deviations to deduce the EUV enhancements associated with flares for which no direct EUV measurements are available, and to correlate them with other flare radiations such as those in X-rays, microwaves, white-light flares and He flares (Donnelly, 1970; Kane and Donnelly. 1971).

5. Conclusions

A. The upper chromospheric line emissions in the extreme ultraviolet show impulsive enhancements in flares which may be statistically related to flare areas in He by the expression:

E = kA 3/2 ,

SOLAR FLARESIN THE EXTREME ULTRAVIOLET 175

where E is the enhancement in percentage of the undisturbed flux, A is the corrected

area in heliocentric square degrees, k is a factor in the range 1 to 3, depending on the

wavelength of the line.

B. These impulsive EUV enhancements peak about 2 rain prior to the He maximum. The EUV rise and decay times average about 2 rain and 5 rain, respectively. There is some evidence in the data that the peak enhancements occur at successively later times for ions formed below, within, and above the chromosphere-corona transition region,

respectively. The ions in the last group, formed at coronal temperatures, also show a much slower rise and fall in their enhancements than do the ions of the upper chromo-

sphere and the transition region.

Acknowledgements

This work was supported, in part, by the National Aeronautics and Space Admini- stration, through the OSO Program Office. I t is a pleasure to acknowledge the cooper-

ation of all the NASA personnel who contributed to the success of the OSO-III satellite. Thanks are also due to Ball Brothers Research Corporation for an excellent

spacecraft, and to Adcole Corporat ion and Comstock and Wescott, Inc. for their technological assistance in the design and construction of the experiment package.

Finally, I wish to thank my colleagues at A F C R L for assistance and fruitful discussion.

References

Athay, R. Grant: 1966, Astrophys. J. 145,784. Castelli, J. P. and Richards, D. W.: 1971, in preparation. Cox, Donald P. and Tucker, Wallace H. : 1969, Astrophys. J. 157, 1157, Donnelly, R. F.: 1967, J. Geophys. Res. 72, 5247. Donnelly, R. F.: 1969, J. Geophys. Res. 74, 1873. Donnelly, R. F. : 1970, ESSA Tech. Report ERL 169-SDL14, U.S. Govt. Printing Office, Washington. Dupree, Andrea K. and Goldberg, Leo: 1967, Solar Phys. 1,229. Hall, L. A. and Hinteregger, H. E. : 1969, in C. de Jager and Z. Svestka (eds.), 'Solar Flares and Space

Research', COSPAR Symp., 81. Hall, L. A. and Hinteregger, H. E. : 1970, J. Geophys. Res. 75, 6959. Hinteregger, H. E. and Hall, L. A. : 1969, SolarPhys. 6, 175. Jordan, Carole: 1969, Monthly Notices Roy. Astron. Soe. 142, 499. Kane, S. R.: 1969, Astrophys. J. 157, L139. Kane, S. R. and Donnelly, R. F. : 1971, Astrophys. J. 164, 151. Kundu, M. R. : 1965, Solar Radio Astronomy, Interscience Pub., New York. Munro, Richard H., Dupree, A. K., and Withbroe, George L.: 1971, Solar Satellite Project, Harvard

College Observatory, TR-23. Neupert, W. M.: 1969, Ann. Rev. Astron. Astrophys. 7, 121. Sawyer, Constance B. : 1967, Astrophys. J. 147, 193. Solar-Geophysical Data: 1967, ESSA, Boulder, Colo., U.S.A.

solar flares in the extreme ultraviolet

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