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X-682-73-21 0 

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SPATIAL DISTRIBUTION OF 

SOFT X-RAY j.\ND ~UV EM.ISSION 

ASSOCIATED WITH A CHROMOS~HERIC I • • 

FLA E OF IMPORTANCE 1B 

ON AUGUST 2, 1972 _ (N ASA-TM-X-6 6299) SPAT I AL DI STRIB UTION OF N73-27720 

SOFT X-RAY AND EUV EMISS ION ASSCCIATED 

iITH A CHROMOSP HERIC FLA RE OF IMPOFTANCE 

1B ON 2 AUGUST 1972 (NAS A) 56 p He Unclas 

$5. 0 0 G3/29 0 99 11 

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'J / ' R. D. CHAPMAN ;. ,.. \ . ' ) -- .... 'foo." ~ • 

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X-682-73-210 

SPATIAL DISTRmUTION OF SOFT X-RAY AN,T) EUV EMISSION 

ASSOCIATED WITH A CHROMOSPHERIC FLARE O:i' IMP ... 'RTANCE lIS' 

ON AUGUST 2, 1972 

w. M. Neupert 

R. J. Thomas 

R. D. Chapman 

July 1973 

GODDARD SPACE FLIGHT CENTER 

Greenbelt, Maryland 
SPATIAL DISTRIBUTION OF SOFT X-RAY AND EUV EMISSION 

ASSOCIATED WITH A CHROMOSPHERIC FLARE OF IMPORTANCE 1B 

ON AUGUST 2, 1972 

W. M. Neupert, R. J. "nlOIDIUI and R. D. Chapman 

Laboratory for Solar Physics 

HASA-Godclard Space FliUht Center 

Greenbelt, Maryland 20771 

~bstract 

Soft X-ray and extreme ultraviolet spectrohe1iographs carried 

b, t he OS0-7 (Orbiting ,Sola r Observator/) have been used to record 

the development of XUV emisaion allsociated with a flare of Importance 

IB on August 2, 1972. Spatial resolution was 20 arc sec and spectral 

re.olution wa. adequate to .elect emission lines or iainating within 

.. ll-defined range. of electron temper~ture between 0.05 x 106 and 

6 30 x 10 It. The data show that heating only the pre-existing coronal material 

&dj .~ent to the flare .ite cannot account for the 80ft X-ray elDi.sion 

.. aaure observed during the event. The flare elDission originating at 

6 ' T ~ 2.3 x 10 It exhibits an impulsive component coincident with an 

e 

~l.ive microwave event. This radiation appears to coincide spatially 

with HO radiation emitted at that time and is centered on the neutral 

liM .eparatina _Inetic fiel d. of oppo.ite polarity. One 80ft X-ray-

.-ittina feature. e.tt.Atad to have an initial electron te.p~rature 

of 2 - 10 x 106 K,for.. durina the ~1.ive pb •• e t..edi&te1y over 

the Ro flare. A .ecODd, ar.ch-llke feature ob.erved at vave1enaths near 

1.91 and e.tiMted to have Te ~ 30 x 106 It is located appr_iMtel, 

35.001 ....... tbe .. eftDt. loth r .. ioU ,haft 1e~th. of about 

1 
27,000 km but transverse dimensions small compared to the spatial 

resolution of the spectroheliograph, i.e., less than 14,000 km. 

The region with highest electron temperature exhibits the greatest 

6 stability in position. This region cools to approximately 10 x 10 K 

in 6 to 12 minutes which is compatible with cooling by conduction to 

the chromosphere. The best association with a dark surge is found in 

an emission line of Fe XIV. 

11 
SPATIAL DISTRIBUTION OF SOFT X-RAY AND EUV EMISSION 

ASSOCIATED WITH A CHROMOSPRERIC FLARE OF IMPORTANCE 18 

ON AUGUST 2, 1972 

w. M. Neupert, R. J. Thomas and R. D. Chapman 

Laboratory for Solar Physics 

RASA-Goddard Space Flight Center Gr ~enbelt, Maryland 20771 

1. Introduction 

Observations from above the earth's atmosphere by means of 

balloons, rockets and satellites provide the only means for observing 

6 direc tly the presence of the hot (T > 5 x 10 K) quasi-thermal plasma e 

frequently associated with solar flares. Early observations with 

broad-band detectors demonstrated not only the existence of short-lived 

X-r,ay enhanc.ements in active regions (OSO-4 data reported by Vaiana 

and Giacconi, 1969) but also suggested that such enhancements may even 

occur in coronal regions not associated with active centers (Zhitnik, 

~rutov, Maljavkin, Mandelstam and Cheremukhin, 1967). Unfortunately these 

early observations had inadequate spatial resolution (of the order of 

1-2 arc min) to provide meaningful information on the shapes of the 

eaitting source" and inadequate spectral resolution to define accurately 

the physical co,"ditions within the regions. Somewhat later, observations 

of solar flare. with improved spatial re.olution but still with broad-hand 

.pectral .endtivity (approximately 2-141.) demonstrated that the sources 

of X-ray eaillilon appeared a. long, filamentary structures about 1 arc 

ain 1001 and 20 arc .ec in dia.eter (Vaiana, Reidy, Zebnpfennig, 

vanapeybroeck ILnd Giacconi, 1969; Be~, Gr1aeva, Mandel' aha, 
2 

Vainstein, and Zhitnik, 1969). Vaiana and Giacconi pointed out that, 

although there were considerable similarities between the shape of the 

H~ flare and the X-ray distribution, there was in addition a br.idge of 

X-ray emission between the two Ha emission centers on opposite sides of 

the magnetic neutral line. In 1969, direct flare observations were made 

in EUV wavelengths (Tousey, 1971) and in hard X-ray wavelengths 

(Takakura, Ohki, Shibuya, Fujii, Matsuoka, Miyamoto , Nishimura, Oda, 

Ogawara, and Ot~, 1971), but no events had yet been observed simultaneously 

in two or more spectral regions. Furthermore, all of these observati ons 

were of short duration so that the evolution of the flare could not be 

studied. 

With the Goddard Space Flight Center experiment on the 050-7 

satellite, we have hed the benefit of prolonged observations so that we 

have been able to view a number of flares from beginning to end. In 

addition, we can observe events simultaneously in four wavelengths 

selected from within the ~avelength region 1.8 to 400~ and by means of 

automatic switching can observe up to ~welve wavelengths in any six 

.tnute period. We·thereby collect information on emitting regions in a 

6 .inale event with electron temperatures ranging from 0.05 x 10 K to 

30 z 106 K. We will first present such observations made during an 

~rtance IB flare on Auaust 2, 1972, and then discuss the t.plicaticns 

of th ••• data OD current flare aadals. 
2. Instrumentation 

The Goddard experiment on 080-7 consists of two X-ray spectroheliographs, covering the wavelength ranges of 1.74 to 7.95Aand 

7.95 to 15.901, and one EUV spectrooeliograph covering the range 120 

3 

to 4001, all of which are co-aligned tc within 5 arc sec. The wavelength 

selection of each instrument is independently controlled by ground command . 

The spatial resolution of the X-ray instruments is produced by multi-grid 

collimators with fields o~ view of approximately 20 by 20 arc sec FWHM, 

and with less than 5% response at 30 arc sec off axis. The X-ray 

spectral bands are defin~d by several balanced filter pairs (Ross, 1926). 

These f ilters are made of materials adjacent in the periodic table, and 

each pair is balanced by adjusting the thicknes~es so that their 

tran8missions at a given wavelength are ident ical except between the 

K (or L) absorption edges of the two materials. One-half of the collimated 

radiation for either X-ray instrument 1.s pauod through one filter and 

the other half through the second filter, each beam being measured 

independently by a proportional counter. TRe difference between detector 

outputs 80 obtained is thus e mea.,tre of the incident energy in a welldefined band of wavelength8. This procedure is the same for both 

X-ray instruments, which differ only in the filter materials and 

proportional counters used. 

The spatial relolution of the EUV spectroheliograph i: provided 

by a Wolter Type II arazing-incidence telelcop (Wolter, 1952), which 

focuses the incident r~diatiOD onto the entrance aperture of the 

lpectro.wtu'. The ltae of the ape ture 11 fixed in ona C!i._ndon to 
a projected angular size of 20 arc sec, but in the other dimension 

it can be selected co be 10, 20, 40, or 60 arc sec. bfter passing 

thro'ugh the entrance aperture, the radiation strikes a grating at a 

o alancing angle of 7.5 and is dispersed to a set of Bendix Magnetic 

4 

Electron Multipliers behind the focal curve of the grating. The spectra.l 

resolution of the system is approximately 11 as determined by the width 

of the exit slits mounted along the Rowland circle of the spectrometer 

directly in front of the detectors. There are six exit slits in all, 

two for each detector. A ma3k in front of the slit assembly insures 

that only one oi each pair is unblocked at a time, however. 

The observations reported here were L~de while the spacecraft 

rastered the instrumental viewing directions over a 5 x 5 arc min area 

of the sun. At the end of each raster, a ~lse was generated which 

stepped the X-ray filter pairs to the next set [see Table 1 for the 

pass bands of the filter pairs], and which switched the EUV mask so 

that the altenlate set of exit slits was unco ered. Unfortunately, 

because of telemetry limitations, only two of the three EUV detectors 

could be monitored. The spectrometer was set so that we recorded 

raster maps at 246.31 and 294.91, alternating with maps made at 219. 01 

aDd 265.01, all with an entrance aperture of 20 x 20 arc aec. 

Since a S x 5 arc min raster takes 61.44 sec to complete, each EUV 

l iDe pair "as obaerved every two minutes, approximately. However, there 

are stE ~ of filters in each of the two X-ray instruaents, so their 

cycle tm. vaa a,.,I'CIIIlDatel, siz __ tes. ('l'be cycle tta. is reduced 
to three minu,tes for the bands from 1.74 to 1.901 and from 14.50 t o 

lS.90A since filter pairs 1 and 4 in each instrument are identical.) 

In addi tion to having X-ray and EUV det ectors, the Goddard 

5 

mrperiment contains a small Ebert-F,.stie spectrometer with a spectral 

resolution of about 11 FWHM for isolating the Ha line of neutral hydrogen. 

It i s fed by the zero ·order ref1eccion from the EUV grating and therefore 

.. aaures the sa radiation from the same location on the solar disk. 

COIIp8rbons with ground-based observations then allow us to determ:lne 

&ccurat'e1y the 8patia1 8ca1e and orientation of our spacecraft data. 

'l1ae final cOlBpOnent of the Goddard 080-7 experiment is an X-ray 

po1arj.eter which is based on the principle of incoherent Thomson 

acatt,ering in a bel;yllium block. The central scattering block is 

aarrounded by six proportional counters, wired together nto three 

pairl" whose relative count-rate •. indicate the degree '.If polarization 

of the incident X-rays. The polari.eter nominally measures X-rays in 

the eMra1 r.oae of rouahly 13 to 30 keV, .1thouah the actual laDle 

Vill.aEJ ...... t de,.Ddi. OD the spectrua of the iDcldent be ... 
6 

3. Data Selection 

In order to evaluate the extent to which the EUV data may be 

contaminated by the appearance of unidentified flare-associated emission 

lines or enhanced backaround at the specified wavelengths, we have 

examined spectra of a large flare on August 7, obtained after flare 

.aimum. 'l1le emtuion line at 246.31 is found to be significantly 

weakened relative to the backaround (being five times more intense 

than the backaround in active region spectra but only one-half the 

intensity of the backaround during the flare). On this basis the 

.pectrohe1iograma obtained at this setting are considered to be unreliable. 

The 265.01 channel is suffiCiently close to the Fe XIV transition at 

,264.801 (Behring, Cohen and Feldman, 1972) to monitor the intensity of 

thi. l ine in active r egions, but it see .. to be blended with other lines 

in the flare spectru., and so is also unreliable for the present study. 

Of the two remaining .pectra1 positions, that at 294.91 correspor.ds 

to no known or observed spectral line either fro. the quiet aun, active 

z.aions or flares. The signal recorded at this poSition is therefore 

,z.au.ab1y due only to .cattered radiation in the spectrometer, which 

a1 .. y. appear. to be about 0.151 of the inten.ity of the nearby, very 

.troaa 303.811ine of He II. Scattered 11aht may aho arise fro. H I 

..t •• ion, .tace it contribute. nearly 90t of the solar photon flux 

_1_ 13001 (Ball and Hinter_er. 1970). the wave1eoath range to which 

Mar 1ithi~nuoride coate. pbotocathode 18 .. uitift. In &111 ca .. , 

_ .. _ dlat tba. .tpa1 at 294.9118 a roup __ " of ..... 1oD 

odatMtiIII ill die _lit _ ..... n. 
7 

1 22 The remaining EUV wavelength at 219.0 monitors the 3s 3p P3/2-

2 2 3s 3d D5/2 transition of Fe XIV. In the flare spectrum the emission 

line at this wavelength has a peak intensity three times greater than 

the level of the underlying continuum, part of which may be continuum 

due to recombination onto He III, although much of it is radiation 

scattered by the grating. It is of intere8t to note that the calculations 

of Blaha (1971) predict that tbe transition at 219.01 should become 

stronger relative 2 2 

to 3s 3p P3/ 2 - 2 2 

38 3d D3/2 at 220.071 as the electron 

density increases from 10 8 cm 

-3 to 1011 cm -3 • Our comparison of active 

region spectra with flare spectra confirms this prediction. 

As a result of the above discussion we consider the interpretation 

of only two of the four EUV wavelensths as betns reliable during flares -

the 294.91 position giving spatial information on emission (principally 

from H I and He II) formed in the upper chromosphere, and 219.01 giving 

the intensity of a density-sensitive Fe XIV line. 

Bes de~ the information from these two EUV wavelengths, we utilized 

a1.ost all of the available X-ray data from our experiment in the study 

reported here. Although the analysis to establish the presence of X-ray 

polarization has not yet been completed, the average output of the 

.po1arimeter's three channels was used to give the time profile of the 

flare's total flux between 13 and 30 keV. Solar radiation is almost 

~ntire1y bre..strah1ung continuum at these energies, but line emission 

bellns to beca.e tmportant at the lower energies measured by the balanced 

fi1t~r technique with our two X-ray spectrohe1iographs. Even though 

Cbeir .pactral reso1utiOD i s not truly .aaochro.atic, each Ros. filter 
8 

or two strong emission lines dominate the detector's response, as 

can be seen in Figure 1 which illustrates the situation for the Long 

X-ray instrument. The ions that produce the dominant emission measured 

by each filter pair have already been listed in Table 1. 

Taking both the line and continuum emission into account, we have 

calculated the incident energy in the spectral bands of several of the 

080-7 filter pairs as a function of T , based on the work of Mewe (1972). e 

The r esultant curves for a constant emission measure of 1050 cm3 are 

8hown in Figure 2. Although. the emission measure actually drops 

rapidly with increasing T in any realistic flare model, each curve 

e 

clearly indicates the temperature range in which the filter pair is 

most sensitive. 

The only X-ray data not directly used in this study were produced 

whenever all filters were removed from the optical path of the 

8pectroheliograph. In such cases, the difference in the count rates 

from the two halve8 of the 8Y8tem provided a te8t of its balance. 
9 

4. Observations 

The selected OSO data taken during the event will be considered in 

two ways. First, we will examine the intensity of the flare emission 

.. a function of tilDe withO'.1t regard to the location of that emission. 

Second, we will consider the spatial distribution of the flare radiation. 

Becauoe this particular flare d~veloped so rapidly only the EUV data, 

with a cycle time of 2 minute., could be used to determine the progress 

of intensity with time. The EUV intensities, given in units of counts 

per 0.125 aec at the point of greatest intensity, are plotted as a 

fUnction of time in Figure 3. A plot of the total flux at each 

wavelength (obtained by summing all enhanced intensities over the 

flare !s area) would be similar except that the flux remains relatively 

h1aber after 1842 UT. The characteristics we observed are: 

1. . A l ow pre flare intensity at the po 8!t!~n where the 

flare ultimately appears and DO de tectable pre flare 

eahanca.ent up to 1837:37 UTe 

2. AD initial rapid brighteniflg, already in progress at 

1838:57 UT, which reaches a mexi.Dum within one minute 

of 1840 UT in the Fe XIV eaission #~d perhaps also in 

the chromospberic eaission. This max~ is coincident with 

an t.pllisive .tcrowave burst reported by the Algonquin 

Radio Observatory of the Nat~onal Research Council of 

Cauda ill the Solar-Geopbysica1 Data Bulletin of the 

IDrina Ental Data Be"lee of .oM. 
3. A secondary maxi mum at 1842 - 1846 UT, bgain in both 

EUV wavelengths, with an enhancement persisting to the 

end of the observing sequence at 1853 UT. 

nte soft X-ray observations, because of t:. .. le longer time between 

repeated observat ions in the same wavelet t h band, are not useable 

10 

for establishing ' the variation of soft X- 1:ay emisshm with time. 

However, the intensity in the band from 1. 74A to 1. ~ A have be n compared 

to the estimated continuum flux at 6.7 keV as a function of time as 

derived from an extrapolation of data supplied by D. Datlowe from the 

University of California at San Diego (UCSD) hard X- ray instrument on 

OSO-7. From this comparison it appears that our dlt a are consistent 

with a rapid rise i n soft X-ray flux from 1839 to 1840 UT , with a 

broad maximum tasting fr an 1840 to 1844 UT. By comparison, an impulsive 

X-ray burst, observed both by the UCSD instrument at photon en rgies 

from 44 to 216 keV and by the Goddard X-ray polarimeter at 13 t o 30 keV, 

was observed to have a sharp maximum at 1839:40 tIT, concurrent wi.th 

the probable maxtmum in the EUV emission. These data are consist ent 

with earlier observations of events at EUV wavelengths (Hall and 

Hinteregger, 1969; Donnelly, 1969; Neupert, 1969) It has been found 

that such impu1siye events are associated with emission lines such as 

o V produced in the transition region (Hall and Hinteregger, 1969; Hall, 

1971. lane and Donnelly, 1971; Wood, Noyes and Reeves, 1972) , or wi th 

radiation formed in the high chromosphere, such as the Lyman continuum 

(Wood, Noyes and Reeves, 1972). Coronal emission lines from high stage. 

of ionization of iron, .uch as Fe XV and Fe XVI, do not in general 
11 

exhibit an impulsive component but become prominent later in the event, 

either concurrently with the soft X-ray peak or even later (Hall ·and 

Hinteregger, 1969; Neupert, White, Gates, SWartz and Young, 1969). 

In the present event, we observe that the rapid enhancement associated 

with the chr~AOsph~ric emission at 294.9A extends to the Fe XIV ion 

formed at T 

e 

6 ~ 2 x 10 K. However, the Fe XIV emission remains high 

after 1843 UT while the chromospheric radiation diminishes. 

Table 2 presents data on the enhanced flux at three of the wavelengd s 

that were observed. Interpolations in the 294.9Adata were based on 

the assumption that 294.9A intensities varied by the same fractional amount 

as the 2l9.0A intensities du ng the impulsive phase . Interpolations 

in the 1.74 - 1.90A band were obtained by fitting i ten.ities derived 

from the UCSD data, extrapolated to 1.871, to our ob.erved values. 

Data in the 13-30 keV band were derived from the average output of 

the three channels in our X-ray polarimeter. We ob.erve the typical 

pattern of maxtDum chromospheric enhancement occurring during the 

impulsive phase and coinciding with the hard X-ray and microwave peak •• 

It is noteworthy that at 1838:38 UT the chromo.pberic brightening wa. 

already at l8~ of i t . maxtDum value whereas the hard X-ray emission 

between 13 and 30 kev has only attained 21 of i t. ultimate inten.ity. 

Maximum enhancement in the 1.74 - 1.901 band,which include. the prominent 

eai.sion line. of Fe XXV, ia reached about one minute after the microwave 

and bard X-ray peak. Maz~ emi •• ion in tbe band containina the re XVII 

liDe. occur. even later. It ba. be.n auaa •• t.d by Heupert .t al. (1969) 

that th. proar ••• lv.ly lat.r .-at.. ln lower .t.... of loniaatlon could 
12 

be attributed t t' .. ucces8ive recombinations in a cooling plasma. Although 

this explanation still seems to be valid. we will be able. after examining 

the spatial distribution characteris ic~ to add significantly more detail 

to such a mode 1. 

Up to this point we have not introduced the spatial f.nfo tion 

contained in the 050-7 data and have not been able to comment on the 

proximity of the emitting regions in space as well as in time. We now 

~ill do so. by considering in turn the spatial emiS8 n of the preflare 

radiation. the emission associated with the impulsive phase and that 

emitted after the impulsive event. 

Figure 4 compares the intensity contours obtaine~ prior to and 

duripg the impulsive phase of the flare. These spectroheliograms are 

constructed by sweepi~ tbe instrument back and f orth over the five arc 

minute field in sixtec .. line. spaced 20 arc sec apart. In these and all 

aubsequent displays of the data. the intensity steps between contours 

are scaled to one-eighth of the maximum intensity recorded in the raster. 

I t il therefore possible to compare spatial distributions taken at 

various wavelenath. and to detect changes in spatial distribution in 

the presence of fluctuations of Intensities. Using this method. we 

note in Figure 4 that at all wavelength. the flare-associated emission 

i. 10 large that the lowelt lev 1 of the :'are cont3Url is in general 

areater than the maxtmua output recorded anywhere in the field of vi ew 

prior to the flare. More f~rtantly. the position at which the flare 

appear I doe. not corre.pond to the loca tion of .axt.u. pref13.6 awi lion. 

either at the EUV V8Velengtba or in t~e .oft X-ray band. containing the 

Mil XI or the re XVII Ii.. (no ai .. ioa whatever va. recorded in the 

bands Dear ~. 91 prior to the flare). Indeed. the only nliao ob.erved 
to be emitting X-rays three minutes prior to the event is associated 

with one of the three r egions observed to be bright ir, H~, and is 

located 1.2 arc mi n (50,000 km) from the flare site. The data have 

not been corrected for the smearing produced by the raster motion of 

13 

the sp~cecraft or the limited spatial response of the system, but such 

corr ections would only serve to emphasize that the flare does not 

originate in any prominent coronal structure existing prior to the flare. 

Repr eseutative observations at xuv wavelengths during the flare 

are shown in Figures 5, 6 and 7. The EUV observations in Figure 5 are 

selected t o demonstrate the evolution in the spatial distribution of the 

r adiation observed during the event. In Figure 6, we compare soft X-ray 

data for t hree selected bands containing the helium-like and hydrogen-like 

emission lines of Ne, Mg and Si (the first two being obtained with th~ 

Long X-r ay spectroheliograph, the third being record~d by the Short 

X-ray spectroheliograph). Finally, in Figure 7, we show a sequence of 

X-ray rasters which include observations of the Fe XXV emission at 1.86k 

During the impulsive phase, which lasts from 1838:30 to approximately 

1842 UT, the flare has a very similar appearance in the radiation of 

He .11 (observed indirectly in Figure 5 by its radiation scattered off the 

grating), Fe XIV (also in Figure 5) and the Mg XI band (Figure 6). 

In each ca8e the source is elongated in a northwest-southeast direction 

with a well-defined bend north of the position of maximum emission. 

We estimate that at this time the source has an overall length of 

36 arc 8ec (27,000 km) and a width that i8 8mall compared to the 8patial 

response of the obaerving ayatema. liIlen we obae.:cve near 1.91, however, 
the region has a distinctly different location and geometry. This 

can best be seen in~igUre 7. First, we must emphasize that when 

both Long and Short X-ray spectroheliographs are recording longer 

14 

wavelength radiation (e.g., at 1840:41 UT when one system is recording 

the band containing the Mg XI and Mg XII lines while the other is 

observing the Si .XIII and Si XIV lines), the observed spatial distributions are similar. However, when tha Short X-ray system observ n.~ ___ _ 

passbands near 1. 91, the spatial distribution is distinctly different., 

even though the distribution recorded at the same time by the Long X-ray 

system has not changed. Such an effect has not been observed in flare 

o events occurring near central meridian (this event was located 30 east 

of the meridian). Neither has any malfunction in the operlltion of the 

X-ray spectroheliographs or instability in the alignment of the collimating 

systems ever been observed. We conclude that this displacement of the 

source emitting 1.9AX-rays is real and amounts to approximately 24 arc sec 

(17,000 km) measured normal to the line of sight. Act.ua11y, this separation 

of e~itting sources is noC unexpected since no single ion can persist 

4 6 

over the large range of electron temperatures (5 x 10 K - 30 x 10 K) 

which are found in this flare, (cf. Section 6). 

Let us consider now the evolution of the flare after the impulsive 

event. Figure 5 shows that after 1842 UT, as the impulsive event was 

coming to an end, the EUV emiSlion at both wavelengths has moved northward 

from its original location, and by 1845 UT the poin~s of max intensity no longer 

coincided. The chromospheric source (294.9]b has elongated and remains . 

r th of its original position until 1848 UT. On the· other hand, the 
15 

Fe XIV source (219.01), while somewhat more extended than·' originally, 

does not exhibit the degree of elongatit)u obserVed at 294.91. Furthermore: by 1845 ~ the maximum intensity of the Fe XIV source has returned' 

to the original position observed at 1839:39 UT. The chromospheric 

source does not return tv Its original location until 1848:52 UT. 

Examining the spectroheliograms taken in the soft X-ray bands 

(Figure 6) , we find that the source may be elongating somewhat more 

slowly than was observed at EUV wavelengths. The location of maximum 

emission remains fixed , at least until 1845 UTe The fact d.at the 

passbands of some of the filters include more than one spectral line 

and continuum emission formed over a wide range of electron temperature 

(see Figure 2) may account for the similarity in the spatial distribution 

of the source in several of the soft X-ray apectral bands. There i. al.o 

some evidence in the data taken after 1851 UT that the center of the 

emission has moved northward and an indication in the 13. 30 - 14.501 and 

7.95 - 9.511ballds that a double maximum may be evolving. 

Considering f1t:.ally the spectroheliograms in the emission lines 

and continuum r. l.9 Which record the spatial distribution of the 

u Ii-thermal plasma having the highest temperature observable with 

the instrument (Figure 7), we find that the position of the source does 

not change throughout the entire event. In addition, its size remains 

always below the r3solution limit of the lpectroheliograph, indicating 

that the hot plasma is lOIDehov·.contained, pol8ibly within a stable 

magnetic confi guration. Thil il consi.tent with the trend toward a 

greater compActne •• with increa.i~, electron temperature already ob.erved 

in the emis.ion line. for.ed at l",rer tUlp8rature. 
5. Comparison with the Associated Ha Flare 

'l11e X-ray and EUV spectroheliograms can be compared with groun<~

based filtergrams using data from the auxiliary Ha system contained 

within the instrument. 'l11e satellite-borne Ha monitor has t he same 

spatial resolution as does the EUV (approximately 20 arc sec), and a 

FWHM passband of about 11. Although these figures are poor relative 

to ground-l,ased equipment, the spectroheliograms so obtained show 

adequate differen(:es in signal from bright plage and from sunspots to 

e8t&blish the spatial relationship between satellite and ground-based 

observations. Because of telemetry restrictions, the OSO-7 Her data 

could not be taken simultaneously with the XUV data. Observations 

11'; 

were taken three hours after the flare, however, and corrected for solar 

rotation to properly register our XUV data with off-b.lnd Ha photographs 

made at the Lockheed Solar Observatory and kindly supplied to us by 

S. Martin. 

The results of these comparisons are given in Figures 8 through 11. 

The spatial distribution of the XUV radiation during the impulsive 

phkse of the ~vent is shown in Figure 8. In all emislion lines produced 

6 at 10 x 10 K or less, the XUV source coincides with the off-band Ha 

distribution to within the accuracy 1!f the spatial cOliaparison, l.e., 

aloNt S arc aec. Fiaure 8a shows this relationship utJ1ng the EUV data 

at 294.91. However, ev n during the iJlpulaive hard X-ray burst, the 

source of radiation at 1.9Ihas a different shape, with ta .ost intense 

point lyina to the north and east of the Ha eaission.. It should be noted 
that the 1.91 spectrohe1iogram scanned through the vicinity of the 

flare within 10 sec of the maximum amplitude of the impulsive hard 

X-ray eminion. 

As the event progresses, the Ha emission elongates tOwar d the 

north, and a break appears in the bt'ight Her filament as it curves 

around a sunspot. A comp rison with magnetograms taken after the 

event by W. Liv sto of Kitt Peak National Observatory i ndicat es 

that t he break i n t he br i ht filament occurs in a region of high 

magnetic field gradients and very near the neutral line. We suggest 

that the no thern and southern positions of the Ha filament may 

therefore be lying on opposite sides of t he neutral line. 

Figures 9 and 10 provide a comparison of XUV emission with Ha 

after the impulsive phase. The chromospheric radiation (294.91) 

still corresponds well to the Ha emission, lthough it may alao bl! 

associated with a dark, northward-moving surge observed in Her. Th~ 

Fe XIV emission maxtDwD, on the other hand, is again centered at its 

original location (after having moved to the north at 1841 UT). An 

elongation to the north does remain and coincides precisely with the 

17 

northward moving surge. The sources of emi uion at 5. 02- 7.951 and 7.95-9.511 show 

1e .. extenaion to the north with ti_. The maxima at theae wavelengths 

coincide with the break in the briaht Her fi1a.nt. The spectrohe1iograma 

at 1.91 show that no motion whatever has occurred in this source fro. 

1839 to 1846 UT j it remains located t o the north and east of the other 

_isS!{'DS and is observed a.ainst a bacqround of low _petie field 

,radiants. Mote that tha eatsdon ppear. at the point where tha dark 

.ursa beco... visibla. 
18 

To summarize these results: 

1. No coronal feature is recorded at the site of the flare 

prior to the event. 

2. During·the impulsive phase, all emissions except at the 

shortest observed wavelength (1. 91) have similar spatial distributions 

and match the most intense region of Ha emission observed off-band 

at 1839 UT. The maximum emission occurs near a region of high gradients 

in the longitudinal magnetic field and, in fact, extendsbetween regions 

of opposite polarity, crossing the neutral line at an oblique angle. 

3. During the post-impulsive phase of the event, the best 

correlation with the structure of the off-band Ha flare' brightening 

i8 found in emi •• i on assumed to originate ' in the high chromosphere. 

Correlation with a dark surge appears to be be8t in the emis8ion of 

Fe XIV. Emission o~1ginating at higher electron temperatures exhibits 

a tendency to remain more compact and stationary. For emission 

6 ' 6 characteristic of 2.5 x 10 K < T < 10 x 10 K, the most intense 

e 

region of the XUV flare i. over a break in the bright Ha filament 

, which may represent the boundary between magnetic fields of opposite 

polarity. 

4. The radiation at 1.91 appears to be di.placed relative to the 

other emitting region.. Moreover, it .eea8 to remain fixed in position • 

and pe.rhap. in .ize, at ~ea.t during the first· 8 minutes of the event 

when iu inteneity 11 adequate for reliable ob.enation •• 
19 

6. PbYlical Properties of the XUV Flare 

Since th lpetial r lolution of the instrument is limited at all 

wav 1e thl and th lpectral relo1ution in the X-ray region is 

iD84equat to 1.~! ~~ individual Ipectral lines, the calculations of 

pbyl1ca1 !)8r at It be only estimates. We will consider 

three lpectral r aionl, (1) l1le v.:ovelength of 219.01 at which the 

2 2 2 2 3. 3p P3/ 2 - 31 3d D5/2 transition of Fe XIV is emitted; (2) the 

spectral band from 7.95 to 9.511 including the resonance transitions 

of Mg XI and Mg XII; and (3) the spectral band from 1.74 to ~.901 

including the transitions of Fe XXV. The relative strengths of line 

emission and continuum in the X-ray bands can be estimated from spectra 

taken by 050-5 on February 27, 1969, during a -flare of imports.nce 2b. 

From Neupert, Kastner, and Swartz (1973) we find that during this event 

the Mg XI "transitions represented 181. of the total emission in the 

7.95 - 9.511 spectral band. For the S8Jte event the ratio of Mg XII to 

6 Mg XI line fluxes corresponded to T • 9.4 x 10 K, based on the 

e 

calculations of Uewe (1972). WE" will adopt this temperature i n OU1r 

analysil. From saba (1972) and Neupert and Swartz (1970) we find 

at the Fe XXV transitions ac~ount for 371. of the total radiation in 

the band bet eu 1.74 and 1. 901. For this rea ion we assume 

T • 30 x 106 K, compatible with the ana1~ is of Saba (1972). r 11 , e 

at 219.01, 751. of the c unt r a te is due to Fe XIV emission, the 

remainder being undisperled radiation IC lt tered by the grating. For 

6 tbil ion we a.lu.. T - 2.3 x 10 K, the t.-perature at which the 

e 

relAtive population of Fe XIV 11 a .-xu... (Jordan, 1969). 
20 

The fluxes in these three spectral regions are given in Table 3 

for the time near 1842 UT, just after the impulsive phase of the flare 

on Auguot 2, 1972. The emission measures listed for Fe XXV and Mg XI 

are based on the calculations of Mewe (1972), Which assume that the 

excited levelo are populated only by collisional excitation from the 

ground state. Average electron densities can then be derived if the 

volume of the emitting region is known. In this case the volumes given 

in Table 3 represent only upper 1tmits. Based on our previous discussion, 

the sources are assumed to be loop-like structures whose length is taken 

to be the distance between half-intensity levels, measured along the 

long dtmension of the source. Because of uncertain projection effects, 

thi& length may be in error by as much as a factor of 2. The observed 

width of each source is equal to that expected when the dimension is 

small compared to the reeo1ution of the collimators. In the absence of 

a reliable value, we will adopt 5 arc sec as the width of the loops in 

these three c.ases. The volume we find for the X-ray and EUV emitting 

26 3 plasma 18 thE!n no greater than 6.7 x 10 cm for any of the three 

emitting regjlons. [It 18 of interest to note that Vaiana and Giacconi 

27 3 (1969) estimated a volume of 1.2 x 10 cm (equivalent to a 15 arc 

.ec cube) for the X-ray emitting pla ... of a flare of importance 1.] 

E.timate. of the electron den.ity and total kinetic energy in the 

la. are al.o liven in Table 3. 

Blahs (1971) has shown tbat the excited levela of Fe XIV lire 

populated by ca.cade. from upper level. a. vell a. by colli.ional 

ezcitation fro. ~e 1r0ua4 .tate. For the 'e XIV radiation, therefore. 

the fiua at the earth 18 liven by 
21 

2 2 where nion,2 is the population of the excited 3s 3d D5/2 level and 

~l is the Einstein rate coefficient to the ground state. This 

expression is equivalent to 

-36 n n n n. -2 F(erg cm -1) _ 7 x ~O J ion,2 ____ i_o.;.;;n__ element hydrogen N 

sec '( nNe A2l dV. ~ ion nelement 'bydrogen e 

thE fractional population in the excited state ni 2/ni is tabulated 

on, on 

by Blaha (1971) as a function of electron density; 

n ion ------- is given by Jordan (1969) and we adopt its maximum value for 

nelement 

Fe XIV (i.e., 0.210) which is reached at T - 2.3 x 106 K; the elemental 

e 

nelement 4 abundance of iron relative to hydrogen we take as 1 x 10-

~ydrogen 

~ydrogen (deBoer, Olthof and Pottasch, 1972); and finally, N we take as 

e 

0.8 in the corona; the gf-value i. given by Blaha (1971) as 2.46 for 

10 -1 this tran.fJition, which means ~l - 5.7 x 10 sec • 

The mean electron den.ity is obtained, for a given volume (again 

.s an upper l~t), by finding that value of ni 2/ni (and hence 

on, on 

electron densit y) which yields a calculated N consistent with the 

e 

initially ••• uaed N. The emission aeasure listed in Table 3 for 

e 

Fe XIV follows fro. this derived density. 
22 

Because of the sequence in which the data were taken, the 

observations used for Table 3 were not simultaneous, but were 

separated by one minute in time. We believe, however, that the 

figures represent r easonable estimates of the physical conditions 

within the quasi-thermal plasma near the end to the impulsive phase 

of the flare. It is true that the flux at 1.91 appears to have remained 

nearly constant from 1840 to 1844 UT, based on our extrapolation of 

X-ray observations obtained by the UCSD hard X-ray monitor on OSO-7, 

and therefore one could argue that the non-thermal impulsive phase 

is still in progress a t 1842 UT, the time of our analysis. On t he 

other hand, if we estimate the e-folding collision loss-time for fast 

electrons interacting with an ambient material of density n as 

o 

8 

(E ) Rf 12- E (kev)3/2 Tcoll e n e sec 

o 

(Trubnikov, 1965; Kane and Anderson, 1970), we find, for E - 100 keV 

e 

and n - 1 x lOll cm- 3, o 

T 11 - 1 sec. co 

Therefore, thermalization proceeds rapidly compared to the 2-3 minute 

d~4tion of the impul.ive burst (implying, however, that there must 

be a continuous injection of electrons during this time). This 

.ugge.t. that the a •• umption of a quasi-thermal plasma may be reasonable 

in our analy.i., even in the case of the 1.91 radiation. 

We can now consider the que.tion of the origin of the X-ray emitting 

pia.... We e.timate from Figure 3 that a pr~flare coronal enhancement 

in Fe XIV at the .ite of the flare could not have been greater than 
57. ~f the maximum flare emission. This would have corresponded to 

46 

an emission measure cf 1.6 x 10 at a nominal coronal temperature 

6 of 2.3 x 10 K. The fact that during the flare emission meas.ures 

6 greater than this were observed at all temperatures up to 30 x 10 K 

demonstrates that the flare emission could not be due simply to the 

23 

heating of such a ~~c -existing coronal structure. The possibility that 

a form of pinch effect is responsible for a compression of the material, 

thus increasing its emission measure, appears to be ruled out by the 

fact that this would require a collapse of the emitting region's 

volume by a factor of 100, which is not observea. We conclude therefore 

that the plasma originated as a result of heating and ionization of 

material origina l ly at near-chromospheric temperatures. Additional 

confirmation of this explanation is found tn Table 2 where we note 

that the inie1al enhancement appears to occur most rapidly in emission 

lines formed at chromospher1c temperatures. 
24 

7. Discussion of the Observations 

We have, up to this time, not discussed the possible spatial 

relationships between the several soft X-ray emissions. The fact 

that the 1.9AX-ray emission does not coincide with any well defined 

photo sph~ric magnetic feature whereas the longer wavelength radiations 

do , suggests to us that the shortest wavelengths originate at a greater 

height above the photosphere. More specifically, we hypothesize that 

the 1.9A radiation may be emitted from an arch-like structure which 

is highest directly above the break in the bright Ha filament, its 

height at this point being about 35,000 km, based on the apparent 

o eastward displacement of the X-ray source and the 30 E longitude of 

the flare. The arch would have its footpoints in or above the 

portions of bright Ha filaments nearest to the break but have little 

or no physical connection with the extremities of the Her. flare. 

To fit ch a picture to the remaining data, we suggest a model 

of the arch structure in which there is a nearly potential magnetic 

field at lower altitudes with lines of force having their planes almost 

normal to ::he ne~Jtral line between the sunspots that are labeled fl 

and P3 by Zirin and Tanaka (1972). At higher altitudes the field 

lines become severely sheared with footpoints near these two sunspots. 

The high temperature source emitting at 1.9A is at the level ~f the 

strongly sheared field, while the radiation of the cooler Fe, Mg, and 

Si sources is formed at the level of the potential field. It should 

be pointed out that Tanaka and Nakagawa (1973) have proposed a model 

.. II 

for the great flare of August 7, 1972, in this same active region, that 

also tmplies that the magnetic field becomes increasingly sheared with 

height. 
25 

Our observations do not permit us to SB.y anything about the 

spatial configuration of material at intermediate temperatures, and 

therefore pres;.unably intermediate heights. On the other hand, we 

believe that there must be a connection betwI!en the two sets of field 

lines to allow for the rapid transfer of energy between them. There 

might be, for instance, a current sheet configuration at intermediate 

heights similar to that considered by Sweet (1958), who a.C'gues that 

magnetic field reconnection takes place in this region, titUS providing 

the source of the flare's energy. Such a configuration wuuld satisfy 

the qui ements of both the XUV and H~ observations insofar as it 

rovid~s a means of quickly dumping energy at chromospheric levels to 

produce short-lived H~ kernels and alsv a source of energetic electrons 

trapped at coronal heights which ultimately heat the evaporated 

6 chromospheric material to temperatures as high as 30 x 10 K. 

To account for the mass motion observed at T < 2 x 106 we must 

e 

assume that only a fraction of the material which is evaporated from 

the v i cinity of the chromosphere or transition region during the initial 

stages of the flare is trapped along the closed magnetic Held ! i nes 

described above. The remainder is eithel' injected with adequate 

kinetic energy to escape from the closed field region or injected into 

a region of open field lines, in either case contributing to the 

observed surge. Such a possibility has already been discussed by 

Gold (1968). 
26 

Using this model, we are now able to discus8 the further evolution 

of the flare plasma after the impulsive phase. We recall that the 

source of 1.91 emission fades rapidly after 1846 UT, so that we have 

no reliable observations of it after that time. We observe in Figure 6, 

however, that by 1852:59 UT the peak of the 5.02-7.95 1 emission 

occupies the site originally emitting at 1.9k Moreover, the emission 

between 7.95 and 9.511 has elongated so that it includes not only 

the original site of that emission but also the region originally 

emitting at 1.91. On the other hand, observations in the band 

containi ng the Ne IX lines show that its source remains fixed in 

space. We propose that he ~ data give direct evidence of the cooli ng 

of the emitting plasma after the UDpulsive phase. Referring again to 

Figure 2 which gives the emission in the bands of seven of our Ross 

50 - 3 filter pairs as ft function of T for an emission measure of 10 cm , e 

we argue t ha t originally just the hottest (and highest) arch system 

can have had a T such that it radiated strongly only at 1.91, i.e., e 

6 it must have had T • 30 x 10 K. The lower arch syste., which vas t: 

observ~d to have nearly the identical geometry in all three spectral 

bands above 51, " lit have had regions with 2 x 10

6 K < T CS; 20 x 106 K. 

6-

By 1852 UT, the upper arch system v.s emitting not at 1.91, but 

in the 5.02 - 7.9S 1band. It auat therefore haVe cooled .0 that 

T < 20 x 106 It. On the other hand, this featun vas not emitting e 

in the He 1][ band, so t t it .aat have had T > 10 x 106 It. The 

e 

upper arch had therefore cooled fro. 30 x 106 to 10-20 x 106 It in 

12 IliDUtea. Alao, by 1852 UT the lower arcb ayat_ vaa no lonaer 
a strong emitter at 5.02 - 7.951. Its maximum temperature could 

6 therefore be no greater than 5 x 10 K after that time. Thus, the 

27 

apparent changes in source location are nothing more than the result 

of chaneing T with time. The ar ch systems th~mBeives have remained 

e 

essentially fixed in space while the maximum T has dropped by a 

e 

factor of -2 and the temperature gradient with height above the 

neutral l ine has lessened. 

These conclusions clearly demonstrate the i ntrinsic limitations 

of spectral observations made without spatial resolution. The 

"average" properties of the flare that are deduced from such data 

may have l ittle resemblance to the actual physical characteristics 

of th~ flare., In instance. where we observe the enhancement and decay 

/ 

in particular spectral line. or bands, for example, the emission from 

one r gion of the flare may be increasing while, at the same time 

the emission from another region is diminishing. 

A feature our picture has in ~n with previou.ly di.cu •• ed 

models (deJager and Kundu, 1963; Takakura and Kai, 1966 ' Sturrock, 

1968; Strau •• and papalianni., 1971) i. the pre. ence of a "Inetic 

f1U1 t ube or loop orilinating in a "Inetic fiel d relion of ODe 

polarity and connecting to a rei ion of oppo.ite polarity. In this 

flux tube are found enera.tic parti cle .tr .... (acc.l. rat.d i n Sturrock'. 

.adel a. the r e.ult of racoDDaction of .. , netic f i.ld 1. n.ut ral' 

.h •• t for.ad by open fi.ld line. abowe th. clo." fi.l. conftauration). 

ADothar property of th ...... 11 11 that the .... rptic particle., 
arriving at the footpoint of the flux tube, produce a hot X-ray 

emitting plasma and may also be responsible for the bright Ho 

points observed simultaneously with the impulsive hard X-ray burst 

(Vorpahl and Zirin, 1970). Although we do indeed observe at least 

one loop system and perhaps several, the region of maximum emission 

(especially for 2.5 x 106 K < T < 10 x 106 K) is greatest not on 

e 

either side of the neutral l ine, i.e., above the bright Ha 

filaments seen later in the flare, but directly above the break 

between those filaments where we 8upp08e the neutral line to be. 

The data are therefore difficult to reconcile with the pr oduction 

of Ha emission by the penetration of fast particle8 into the 

chromosphere as had been suggested by Vorpahl and Zirin (1970). 

A stmilar conclusion has already been made by Takakura (1971) on 

the basis of hard X-ray observations of a flare made by Takakura 

and co-workers (1971). An alternative possibility, of course, is 

simply that our instrument has inadequate spatial resolution to 

establi8h the relatiunship between soft X-ray and Ha emission. 

28 

One fact ~hich the data clearly establish, however, i8 that the 

highest electron temperature is not found (even during the impulsive 

pha8e) immediately above the chromo8phere, where the fast electrons 

are pres~d to be heating the material. The location of the 1.91 

source requlree that the 1ectl~n te~rature increase fro about 

6 6 5-10 x 10 K directly above the chromosphere to 30 z 10 K at a 

he1sht of 35,000 ta. Thus, heat ina and ioni.atlon .ust occur over 

this ranae of heiahts. Such a require.ent does not, however, rule 
I 

out the possibility that non-tnerw&l bremsstrahlung is produced 

es the result of deceleration of electrons in the chromosphere. 

A dynaolic model of the event would have to account for an initial 

~ 

heating of material, perhaps produced by collisional losses of high 

energy electrons in the chromosphere (Neupert, 1968), followed by 

additional heating at heights of 10,000 km or more above the 

chromosphere. Such heating could be produced by high energy 

particletl trapped in .1 closed magnetic field configuration which 

as yet have not impacted on the chromosphere or escaped into the 

high corOlla. 11le entire range of heights over which this heating 

must occur is not ~bservable at any single wavelength since it 

encompaese tl a wide rangtl of eleoctron temperatures. From an 

inspection of our present data we conclude, however, that at 1839 UT 

it eletends from the 294. ~ .1 source observed in the high chromosphere 

(Figure 5) to the source observed at 1~9.1 at which the highest 

electron temperature is to be found (Figure 7). 

Strauss s.nd Pagagiannis (1971) have used the model of Sturrock 

and Smith (1968) and Sturrock (1968) to calculate the expected 

densities and temperatures in a loop when the injection of energetic 

particles occurs at t he top of the loop. Using an injection rate 

which is chosen to give consiatent results at the beginning of the 

event, they are .lble to match the observations for the post-maximum 

29 

pbafJe of an X-ray buret recorded by OSO-5. 'l11ey are beat able to fit 

the data with a loop height of 57,000 ka, maximum T 

e 

of 61. 7 x 106 K, 

. 11-3 and N - 4 x 10 em • 

e 

These conclusions are based on a cros.-

sectional loop area of 7.5 x 1014 ca

2

• H this were to be increaoed 
30 

d ( 1 16 .l to the dimension we have use 9 x 0 cm) and the loop height 

decreased by a factor 2, we estimate that the calculated electron 

density would be only a factor of 3 higher than what we have reported 

earlier in thi s paper. However, a difficulty remains in t hat we find 

no supporting evidence fo r the hypothe sis that energetic particles 

are injected at the top of our loop and thermalized a t the footpoint.,. 

A similar ana1yds has recently been carried out by St urrock (1972). 

His results for a small intense flare having a hot plasma loop 

30,000 km long al;e also in general agreement with our observations. 

Culhane, Vesecky and Phillips (1970) have considered the cooling 

mechanisms for 10()ps such as we have observed. Although they were able 

to rule out collisions of electrons with lower temperature ions as 

the cooling mechanism, they were not able to discrtminate between 

cog1 ing by conduction and by radiation on the basis of available data. 

In our case, however, the fact that the loop at 1.91 is observed for 

6 

more than 6 minutes (i.e., T > 10 x 10 K for this duration) places e 

constraints on any possible mode l. . Since the radiation cooling rate 

-23 

per unit emission measure i8 nearly con8tant at a value of 10 

erg sec1 for temperature8 between 10 x 106 and 30 x 106 K, we may 

write 

or 

dE _ 3 N V k dT _10- 23 N2 Vera 8ec-1 

dt e dt e 

bt. 3k AT 

10-23 R 

e 
31 

Using ~T = 20 x 106 K and N = 1.5 x 1010 cm- 3 we find ~t = 5 x 104 sec, e 

which is more than a factor of 100 longer than the observed duration. 

N would therefore have to be greater than 2 x 1012 cm3 if the rate 

e 

of cooling were due only to radiation. As this density is unrealistically 

high, we feel that radiation cooling is not important in the present 

instance. 

Consider now the possibility of conduction cocling. Culhane et al. 

(1970) discuss conduction models for several values of a parameter a, 

where a ~ 2 x 1010 N -1 t- 2 , and t is the length of the arch in cm. 

e 

9 10 -3 Usi.ng J, ~ 6.4 x 10 cm and N = 1.5 x 10 cm, we find 

e 

-20 

a - 3.2 x 10 • For this value of a the graphs presented by 

Culhane et ale (1970) show that conduction would cauee Te to drop to 

10 x 106Kin 90 sec, i.e., more rapidly than observed. An increase in 

Ne would Qnly produce a greater discrepancy. We therefore conclude 

that the hot plasma may not be coupled directly to the chromosphere 

but may rather be contained in a partially leaky magnetic bottle 

which permits a conduction cooling rate about 4 times less than 

otherwise would be the case. Under such circumstances we expect 

that radiation cooling could still have a minor role. 

In addition to the direct observations of loop-like or 

filamentary structures emitting soft X-rays (Vaiana et a1 :,1968; 

Beigman et a1., 1969), indirect evidence for an arch-like structure 

emitting .oft X-rays a. well as Ha radiation has been presented by 

Zirin, Ingham, Hud~on and MCKenzie (1969). Their argument i. based 

on the observation that the flux .a.ured in the 7.7 - 12.6 keV 
32 

8pectral region by the UCSD in8trument on OSO-3 is well correlated 

with the projected area of a risi ng ~ arch during the increauing 

pha8e of an X- ray burst. Although Zirin et ale reported a maximum 

height of 46, 000 km for their arch, a figure comparable to our results, 

there exist significant differences which make it unlikely that our 

loop at 1. 91 can be attributed to the same type of feature observed 

by them. In the first place, their H~ feature rose at a nearly 

constant velocity (we estimate a component normal to the line of 

-1 sight of 136 km sec ). The feature we report remains fixed throughout 

the observing interval although motion would easily have been de tected. 

Secondly, Zirin et ale concluded that the radiating electrons were 

uniformly ~istributed throughout the arch (because of the smooth 

l ' h t l with time in the X-ray flux). Our feature is definitely not 

uniform in intensity - one side i8 obviously more intense than the 

other. I f it were to ri8e from behind the limb at the rate observed 

by Zirin et al., it would reach nearly maxtmum visibility in 80ft X-ray 

emt88io~ in ~bout 25 8ec. Such a sharp increase was nut observed in 

the OSo-3 data. 

Efforts to define the region in which electron acceleration 

takes place have as yet produced no conclu8ive results. Zirin, Pru8S 

and Vorpahl (1971) argue, from the apparent frequency cutoff in radio 

e.tssion during one event,that the process must take place at densities 

11 -3 of 10 em or at an initial teaperature of 30,000 K. Kane and Lin 

(1972) on the other hand aque, fro. the detection of low eneray 

interplanetary elec~roD.. that the acceleration occurs at 
We are not able to identify any coronal 

structures prior to the flare that might provide the site for 

electron acceleration. A possible solution to the dilemma may be 

in our observations that some field regions are closed and 

effectively trap the high t emperature plasma while othe4s are 

sufficiently open to permit the ejection of material. Thus, it 

might bl~ possible for energetic electrons t o "leak" around an 

overlying cloud of high density plasma which provides adequate 

attenuation for the line of sight radio emission. 

33 
34 

8 • S uDlDllry 

On the basis of spectroheliograms at selected wavelengths 

throughout the XUV region, we come to the following conclusions. 

1. There is no detectable coronal feature at the site of the 

flare prior to the event. Heating only the pre-existing coronal material cannot 

account for the flare emission. 

2. 6 Flare emission at T ~ 2.3 x 10 K exhibits an impulsive e 

component which has already reached 18% of its ultimate intensity at 

1838:38 UT when X-rays at wavelengths below 2Aare just beginning to 

be emitted. This suggests that the flare has its origin in the 

chromosphere or transition region and that energy losses by high-energy 

electrons may not be cessary for the production of the impulsive 

chromo spheric radiation. 

3. With the exception of the highest temperature plasma, the 

source of the radiation associated with the tmpulsive phase follows 

the neutral line between fields of opposite magnetic polarity. The 

emission is centered at a break in the bright Ha flare i lament ch 

probably separates the two magnetic field regions. There is no 

close correspondence between the spatial distributions of the Ha 

briahtening and soft X-ray emission after the impulsive phase, however. 

4. At least two distinct structures are observed at X-ray 

wavelengths. PrOIa the Wavelengths at which emission occurs, we infer 

6 

an initial teaperature of 2-10 x 10 K for one structure which forma 

t..ediately over the Ba -flare during the t.pllisive phase. Siaultaneousl" 

6 

an arch with T A$ 30 x 10 K develop. about 35,000 1aa above the sa 

e 

event. The reliona with htabe.t electron ta.perature exhibit the 
greatest st ability in p~sition and perhaps in size. We conclude 

that the magnetic field l ines directly above the principal source 

of energy dissipation are closed. The observations require taat 

ionization and heating of evaporated chromosphe ric material continue s 

as the material streams upward. It may therefore be that t he region 

of 1.91 emission de1ine&tes a storage region and perhaps even an 

acceleration region for high energy e lectrons. As the particle 

density increases in this region during the impulsive phase, the 

collisional energy loss rat e will increase and may "quench" the 

electron accelera tion process. 

35 

5. From the long lifetime of the high t empe r ature region (at least 

6 6 minutes and perhaps as much as 12 minutes to cool from 30 x 10 to 

6 10 x 10 K), we conclude that this r egion may be conductively decoupled 

from the chromosphere. Even so, radiation losses are significantly 

smalle r than conduction losses. 

6. In addition to the stationary plasma at high electron 

temperature, we al so observe moving material apparently associated 

with a dark Ha surge. Thi s is best observed in an emission line of 

Fe XIV during the post Mimpulsive phase of the flare. The ejection of 

this material occurs not from the top of the closed a rches but either 

from below them (and therefore directly at the neutral line) or from 

the vicinity of the northern foo tpcints of the arches. 

It has become clear to us that the instrument used for these 

ob.ervations is barely adequat e to tree the evolution of tnis 

ca.pl ex event. However, even these rudt.entury results live .. ple 

evidence of the rich fields of observational and theoretical results 

whi ch r ... 1n to be uodertaken in the physics of solar flares. 
36 

Acknowledgements 

Many people have contributed to the success of our program of 

observations from the OSO-7 satellite. We are greatly indebted for 

the competent technical management provided to us by Milton Kalet and 

Gary Harris of the Goddard Space Flight Center and for the contributions 

of our co-investigator, Dr. James H. Underwood. now at the Aerospace 

Corporation. Assistance in calibrating and operating the experiment 

has been provided by Charles Condor, Edward Rothe, Robert McColm, 

Marvin Swartz and Robert Young. The instrument was built for us by 

the General Electric Co., Valley Forge, Pennsylvania. We wish 

part1cularly to thank Andrew Sterk, Frank Kieser, Earl Knox and Stephen 

Peck for solving many technical problems and for providing an instrument 

capable of achieving our scientific objectives. The spacecraft was 

built by the Ball Brothers Research Corporation and has proven to be 

a reliable and stable platform for our observations. We wish to 

recognize the efforts of the OSO froj ect Office and the OSO Control 

Center at Goddard in managing the preparatio of the spacecraft before 

launch and providing dependable operations thereaft r. Their assistance 

has been indispensible and is sincerely appreciated. We al 0 wish to 

ackngwledge several informative discussions with Y. Nakagawa on the 

configurations of magnetic fields in active regions and flares. All 

Ha photographs in this paper were graciously supplied by S. Martin 

of Lockheed Solar Observatory. One of us (WHN) would like to express 

his thank. to Dr. Gordon Newkirk, Director of the High Altitude Observatory, 

for his hospitality during a sabbatical leave when work on this paper 

• ~a.pleted. 
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37 
38 

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39 
Table 1 

Short X-Ray Spectrohe1iograph 

Filter • ~(A> Dominant PO'l.tion Element A( A) Emission Lines 

1 Fe - Mn 1.74 - 1. 90 0.16 Fe XXIV, Fe XXV 

2 Mn - Cr 1. 90 - 2.07 0.17 Fe sate1U tea 

3 S - Al 5.02 - 7.95 2.93 Si XIII, Si XIV 

4 Fe - Mn 1. 74 - 1. 90 0.16 Fe XXIV, Fe XXV 

5 Fe - Cr 1. 74 - 2.07 0.33 Fe XXIV, Fe XXV 

6 (None) 

Long X-Ray Spectrohe1iograph 

1 Ni - Co 14.50 - 15.90 1.40 Fe XVII 

2 Ge - Zn 10.19 - 12.13 1.94 Fe XVII, Fe XXII, Fe XXIV 

3 (None) 

4 Nt - Co 14.50 - 15.90 1.40 Fe XVU ~ 

0 

5 Cu - Ni 13.30 - 14.50 1.20 Fe XVII; Ne IX 

6 Al - Mg 7.95 - 9.51 1.56 Mg XI, Mg XII 
Table 2 

- -- -" -

Tt. t 30 He 

Flux 

(Total counts ler 0.1:5 lee) (UT) 294.91 14.5O-1S.901 1.74-1.901 13-JOkeV 

1838:38 994 162 Q2 13 

1139:39 ~S80 (eat. ) 16S4 28S (eat. ) 664 

,I 1840:41 4268 2454 (eat.) 300 (eat.) 38S 

1141:42 ~7 (eat.) 2626 284 206 

1847:50 1728 1850 68 A() , 

aiL ,-

Percent of 

MaxiDllD Flux 

294.9.6. 14.50-1S~90A 1.74-l.90A 

18 6 ~ 

100 63 I 9S (est.) 

76 93 (eat.) 100 . SS (eat.) 100 9S 

31 70 23 

l3-30keV 

2 

100 

58 

31 

A() 

-

~ 

.... 
Table 3 

Characteristics of XUV Flare near 1842 UT 

..... 

T F N

2 V V 

Ion e -2 -1 e_3 3 (106 K) (erg cm sec ) (cm ) (cm ) 

Fe XIV 2.3 1.3 x 10 -3 3.2 x 1047 60 . 7 x 1026 

_'XI 9.4 1.4 x 10 -4 3.9 x 1048 ~.1 x 1026 

Fe XXV 30 .•• 8 x 10 - 6 7.2 x 1046 ~.3 x 1026 

N 

~3 (cm ) 

~.2 x 1010 

~1.1 x lOll 

~1.5 x 1010 

3 N V kT 

e 

(erg) 

1.4 x 1028 

1.3 x 1029 

6.1 x 1028 

~ 

~ 
20lI0' 

I I 

11 III 

• • 

.. 

f 

I 

" 

• f 

• 

I 

'f 

l 0 I 

... 

y 

~ 

/ -

01 III d ql 

-.uDIGfH 

43 

i 

I 

n 

\ 

Figure 1. Soft X-ray transmission bands for the Ross fUter pairs used in the 

Long X-ray Spectrohellograph 00 080-7. For illustrative purposes the figure 

includes the X-r~ spectrum C!'I a 2b flare ntwt~f!led 00 FebnlB.l'y 27, 1969, by 

the Goddard 080-5 experiment. In moet cases, only one or two strong 

emis8ion lines dominate the bandpass of each fUter pair. 
1.0 

.., I 

e 

u 

St 

0 

-II lxl0- > 1 o Ne 

Z • • 5.02·7.95 A 

! 

~ 

--I 

= N Ie 

u 

If, 

.. 

~ 

u 

.., 

+ 

..I .., 

" o 

_ ---- 1.74-1.90 A ). .-=- 0 < 0 .-.",.,.- tv 7.95-9.51 A 

/ 

/ 

/ 

/ 

" 13.3-14.5 A 

I' I / I I I 

l

x

l0-

3

0 10 20 30 40 

Te (106 K) 

50 

Figure 2. Line aDd conUnu,.1Dl emission as a function of electron temperaturE" integrated 

over the apectral range of several Ross filter pairs used in the 080-7 instrument. The 

incident energy in each band is calculated for a constant emission measure of 1050 cm-3 • ~ 

~ 
45 

~r---------------------------------------------------------,I~ 

u 

.... 

(/) 

.... 

N 

o 

700 

600 

500 

III: 400 

~ 

(/) 

~ 

Z 

:::I 

o 

h 

, \ 

I \ 

, \ 

I \ 

I \ 

, \ 

I \ 

I ' . \ \ I I \ \ . , I \ \ 219.0 A 

I I \ \\ FE XIV EMISSION: 75' 

\ CONTINUUM + RADIATION 

I I \ \ 

1400 

1200 

1000 

BOO 

. u 300 294 .9 ;. 

DATA 

I I \ \ SCATTERED1BY GRATING: 25' 

, I \ \ 

I I ~ 

219.0 A 

DATA 600 

I I ~, / .... 

200 

, I ' ........... ~ ........ .-00. ........ , '" -II'" _- -, .... _ _----

, I ~,~--- - ~ 

/' 

If 

294.9 A /" '0....._ 

100 /' CONTINUUM . RADIATION --0... ___ " SCATTERED BY GRATING 100'1 ---"""'0 ---.'1 

-&>-_-cr-J 

200 

1834 1836 1838 1840 1842 1844 1846 1848 1850 1852 

UNIVERSAL TIME 

Figure 3. The intensity of the brightest point in the EUV flare of August 2. 1972. 

at two wavelengths: 219.oA and 294. 9A. Note the impulsive phase between 

1838 and 1842 UT and the gradual phase after 1842 UT seen at both wavelengths. 
183606 - 1837·07 

294.9A ICHROMOSPHERIC RADIATION) 1838 OS - 1839 10 

", . 

. . .. ... .. " 

" . :' ... .. , ' . ' '. 

1837.07 - 1.38 oe , .. . '. ... >, • 

. -', 

" , . . 

219.0A IF. XIV) 1839-10 - 11140-11 

", ':. ,:. . :.,', . ~": ' .. : . ',--' 

1 .. :01 - 1":01 

" :'" ... " 

14.50.15.90A , .. XVIII 

.. 

r 

: ,' 

FilUre 4. SpeotroheUograms of solar emission 

prior to and during the impulsive phase of the flare 

at two EUV wavelengths and one spectral band of soft 

X-rays. All spectroheliograms are 4 arc min on a 

side and are oriented with north at the top and east 

to the left. Universal time at the beginning and end 

of each raster map is shown. Intensity levels have 

beeD normalized so that, in each spectr ohellogram. 

the iDcrement between two lsophotes represents 

one-e1gbth of the maximum lntensity found in that 

apectroheUocram. Note that only weak EUV and no 

measurable soft X-r&Y. emission emanated from the 

site of the flare before the event bepn at 1838 UT . 

Preflare emls.10ll1. u80Ciated with repone that 

are brtpt 1n Her • 

46 
. . ; ...... ::.: ~., 

;:;t~~:: .. ". 

... ., 

.. ' .... " ... .... '. •. 

\::=:*~~f~) 

,",Q 

: " \:. 

;::&o~.'. \ , .... ,. .d .. . ::1 ;;~::' 

2M .• 1 lCMIOIIO ..... te IADtAnoNl 

219.oA IF. XlVI 

Figure 5. Spatial evolution of EUV flare sources. The earliest frames show the location of the impulsive 

event at about 1839 UT • Thereafter. the sources become elongated toward the north. but eventually (after 

1049) tend to contract again to the initial site of the event. The times indicated on this and following 

figures refer to the actual time at which the raster scan passed through the center of the flare. 

~ 

...::I 
13.30 -14.50 l 

",IX+F. XVII) 

I ••• 

7.t5-9.51l 

(IIa XII ..... Mill) 

5.o2-7.t5l 

(liXIY • liXIII) 

1840:41 

I14U7 1151:56 

I14U' 

, 

Figure 6. Flare emission in three solt X-ray bands containing emission lines of ions 

in the H I and He I isoalectronic eequences. Ions responsible fo~ the major emission 

lines within each band ara indicated. The upper two rows of data were acquired with 

the Long X-ray instrument and the third row with the Short X-ray instrument. Note 

the elongation of the souroe betweElll 1839 and 1846 UT. By 1852 there is a definite 

indication that emission in the 5.02-7. 95A band has shifted toward the location 

previously occu.ried by the 1. 74-2. 07A source (see next figure). ~ 

00 
1839:39 

13.3 -14.5A 

II 

1839:39 

o 

1.90-2.07A 

1840:41 

1840:41 

5.02-7.95A 

1842:43 

10.19-12.13A 

1842:43 

1.74-2.07A 

1846:49 

7.95-9.51A 

1846:49 

5.02 -7.95A 

Figure 7. Comparison of soft X-ray spectroheUograms taken cODCurrently with the Long X-ray and Short 

X-ray SpectroheUographs. Note that whenever radiation near 1. 9A is being recorded. the source is 

displaced to the north-east (apper left). ~ 

to 
(c) 

(b) 

Ha + 0.61 AT 1839:46 UT 

(o) 294.91 AT 1838:38 U 

(b) 7.95- 9.511 AT 1840:" UT 

(c) 1.90-2.071 AT 183~ 39 UT 

50 

Figure 8. Comparison of XUV spectroheliograms ta..'rea during the impulsive 

phase of the flare with an off-band (Her + 0.6 1..) photograph made at the 

Lockheed Solar Observatory . 
Ha + 0.61 ;r 1847:30 UT 

(0) 294.91 AT 1846=49 Uf 

(til 7.96-9.511141 1846= UT • 

(d 190-2.071141 1845~47 UT 

51 

Figure 9. Comparison of the XUV spatial distribution during t1.le post-impulsive 

phase of the flare with an off-band (Hex + o. 6A) photograph made at the Lockheed Solar Observatory . The XUV wavelengths are the same as those shown 

in Figure 8. Note that the radiation at 294. 9A attributed to chromo spheric 

emission coincides with the Hex brightening present at nearly the same time. 
52 

Figure 10. Comparison of an EUV spectroheliogram in the radiation of Fe XIV 

with an off-band Her photograph made at the Lockheed Solar Observatory. This 

ion exhibits the best association with a dark northward-moving surge as shown 

here at 1847:30 UT. The 219. OAmap was made at 1845:47 UT. 
53 

Figure 11. Superposition of an EUV spectroheliogram at 294.9A taken during 

the flare's impulsive phase at 1838:38 UT and an off-band Ha photograph recorded as the flare was subsiding at 1847:30 UT. Note that the impulsive 

EUV radiation originated directly over a dark gap in the bright Hex filament . 

This gap coincides with the neutral line which separates areas of opposite 

magnetic polarity in the active region .