temporal trends of solar euv and uv full-disk fluxes

T E M P O R A L T R E N D S OF SOLAR EUV AND UV F U L L - D I S K

F L U X E S

R I C H A R D F. D O N N E L L Y

Air Resources Laboratory, NOAA ERL, Boulder, CO 80303, U.S.A.

(Received 12 November, 1986; in revised form 5 February, 1987)

Abstract. Several progressions in the temporal characteristics of full-disk solar UV and EUV fluxes have been identified that raise many questions about the solar physics involved. The collective effect of numerous enhancements smaller than scaled plages contribute significantly to the solar cycle variations, especially for emissions from the cooler portions of the corona and the chromosphere. Active-region remnants are suggested to have a strong role even in solar-rotation induced variations late in an episode of major activity. Although cool coronal EUV emissions are long lasting, the persistence of the solar-rotation induced variations is even greater at photospheric UV wavelengths. Gyroresonance and possibly nonthermal radio emission at centimeter wavelengths are suggested to be particularly important during the first solar rotation of an episode of major activity.

1. The Solar-Terrestrial Photon Connection

The temporal characteristics of full-disk solar extreme ultraviolet (EUV) and UV flux measurements have recently been studied primarily for research of their terrestrial atmospheric effects (Donnelly etaL, 1986a, b; Heath and Schlesinger, 1986; Hinteregger, 1986; Rottman, 1985; Schmidtke, 1984). However, some of the results raise questions about the physics of the solar causes of the observed temporal variations that should be of interest to solar physicists. This paper discusses some of these observed temporal characteristics and emphasizes the solar physics questions raised.

Since OSO IV measurements of EUV spectroheliograms in 1967 (Reeves and Parkinson, 1970), EUV and UV observations for solar physics research have tended toward higher spatial, wavelength and temporal resolution, which has led to an increasing gap between solar research and studies of full-disk flux measurements for research of the terrestrial atmosphere, where the latter emphasizes absolute accuracy, completeness of wavelength coverage and longterm temporal variations. This paper attempts to reduce that gap by raising questions about the longer-term evolution of solar activity that may so far have been missed in the high resolution solar physics data.

2. EUV and UV Measurements

The EUV flux data studied were measured from the AE-E satellite and are the same as those studied by Donnelly et aL (1986a, b), namely the 15 wavelength groups of flux ratios F/FRE F with the file name SC #21OBS (Hinteregger et aL, 1981), where Fis the measured flux for a particular day and wavelength group. The subscript REF denotes the AE-E reference period, 13-28 July, 1976. FRE v is the absolute flux for a particular wavelength group calculated from the reference EUV spectrum named SC # REFW

Solar Physics 109 (1987) 37-58. �9 1987 by D. Reidel Publishing Company

38 RICHARD F. DONNELLY

(Hinteregger et al., 1981), which lists estimated values of irradiance for essentially all known solar EUV emissions and are recommended as the AE-E solar-minimum values for solar cycle 21. The SC #OBS data set starts on 1 July, 1977, during the early rise of solar cycle 21 about one year after the sunspot number minimum, continues through the main rise to the peak of the sunspot cycle (end of 1979 and start of 1980) and ends on 30 December, 1980. Table I describes the EUV data studied, including the source regions involved.

Bossy and Nicolet (1981), Bossy (1983), and Oster (1983) have criticized the AE-E measurements as suffering from shifts in the instrument sensitivity to EUV radiation, especially the H L~ measurements. Except for possible long-term changes in the AE-E H L~ measurements, Donnelly et al. (1986a, b) disputed these criticisms. The AE-E EUV measurements were made by a group ofmonochromators. The L~ data were based on measurements made by a fixed-wavelength monochromator (MN 22) that was essentially separate from the wavelength-scanning monochromators providing the observations for the remaining fourteen wavelength groups. If a fault occurred in this one monochromator, it is not likely that a similar fault would occur in another quite different and independent unit at the same time. Furthermore, Donnelly et al. (1986a, b) showed that the differences in the temporal variations of the chromospheric EUV fluxes with respect to the solar 10.7 cm flux (FI0) were like the concurrent differences with respect to F10 shown by independent NIMBUS-7 measurements of photospheric and chromospheric UV fluxes and ground-based measurements of the chromospheric He I absorption line at 10830 A.

The UV data studied were measured by the SBUV experiment on the NIMBUS-7 satellite in the 1600-4000 A range with an 11 ,~ half-max bandwidth on about three out of four days (Heath and Schlesinger, 1986). The flux at 2050 ,~, a wavelength in the Ah absorption continuum that is slightly shorter than the A11 absorption edge, is used here as a representative case for the NIMBUS-7 UV flux measurements. The 1-8 ,~ soft X-ray flux data studied were daily means and the midday background flux from the GOES-2 satellite, with occasional cases of missing data filled with data from the GOES-3 satellite (Bouwer et al., 1982). These 1-8 .~ measurements are very sensitive to the coronal emission measure at temperatures T > 3 x 106 K and insensitive to the emission measure below 2 x 106 K (Donnelly et al., 1977). The ground-based measurements of the full-disk equivalent width (EW) of the HeI absorption line at 10 830 A were discussed by Harvey (1984). Although spatially resolved HeI data show a sensitivity to coronal radiation, the temporal variations of the full-disk data are dominated by the chromospheric fluxes of active regions and shown to be highly correlated (r = 0.97, Harvey, 1981) with the full-disk Call K-line measurements of White and Livingston (1981). The sunspot number (R), Ottawa 10.7 cm solar radio flux (F10) and the daily Ca-K plage index (P) were obtained from Solar Geophysical Data.

3. Temporal Variations

Figure l(a) shows four examples of Fe lines observed by the AE-E monochromators. Three well-known temporal variations are evident: (1) the short-term variations (days,

T E M P O R A L T R E N D S O F S O L A R E U V A N D U V F U L L - D I S K F L U X E S 39

14 o~ >; ~5 12

~ 1o n -~

o ~bx 8

c

o

:i "J ttl' ,,,,

2 I

1977 1978 1979 1980

~o 2.4

>~ 2.2

2.0 • L~ "5 1.8 ' sx ~_ 1.0

~ 1.4

g~ 1.9

~, 1.0

AE9E17MN~ 1

Domina led by Fe IX IIl~/~[~iI/lil ~~ i i

1977 1978 1979 1980

AE*E MN 1 ~ ] t t l tt ~ t ~ 178-183,8, Domina ted by Fe Xl , i l l ~I

~ 1977 1978 1979 1980

Fig. la. Temporal variations of EUV Fe lines. MN refers to the monochromator number on the A E - E satellite.

weeks), induced mainly by the solar rotation of an inhomogeneous distribution of active regions with respect to solar longitude, (2) intermediate-term variations (typically 4 to 8 months), seen in the slow rise and fall of the rotational minima and also in the rise and fall of the valley-to-peak amplitudes of the rotational modulation, which are asso-

i

T.

'5

Fig. lb. based

tO 350 m 3.5 AE-E MN 5

>; 584 ~, He I 300 ff~

3.0

250 -- L_~ O 2.5 I o x , lY 2OO o--~ l[ t L =LL

o~ 2.0 !1/}~1 ~50 ~ i ~

1.5 ~1 100 o

~" 1,0 1977 1978 1979 1980 m" 1977 1978

7qq- rqqq 'T I I I L I I I ~ I I , ~ , I I I I I h I I ~ I ~

Ca-K Plage Index P

L~JJJ~ I I I I',,, 1977 1978 1979 1980

1979 1980

E

Temporal variations of a He I EUV emission line measured from the A E - E satellite and ground- measures of solar activity, including the equivalent width of an infrared HeI absorption line.


ciated with the evolution of major groups of active regions, and (3) the rise and early maxima of solar cycle 21. No connecting line is shown when more than four consecutive days of data are missing. The lack of connecting lines shows clearly that the FexIn measurements are available for fewer days than the FexvI measurements and that more data is missing in 1977 and 1978 than in 1979 and 1980. Consequently, the short-term variations, discussed in Sections 3-7, were analyzed for 1979-1980. Note that the size of the short-term variation relative to the solar cycle increase varies among these four Fe lines. Fe • observations at 284 A (not shown) have temporal variations very similar to those fo FexvI, except the dynamic range for Fexv is about one third that for the FexvI observations (r = 0.98, Donnelly et al., 1986a).

The temporal characteristics of the chromospheric HeI 584 .~ line are quite similar to those of the 10 830 .~ line in Figure l(b) (r -- 0.90) and very similar to the AE-E observations of the H Lfl line (not shown, r = 0.98, Donnelly et aL, 1986a). The 10.7 cm flux (F10) corrected to 1 AU looks most like the hot coronal FexvI observations, including large short-term variations relative to the solar cycle increase and a sharp spike in November 1979. The amplitude of the short-term variations of the chromospheric Ca-K plage index (P) are large relative to its solar cycle increase in Figure l(b), much more so than in the cases of the two independent HeI observations. This suggests that the scaled Ca-K plage data are insufficient to explain the long-term variations of chromospheric fluxes, as has been shown in the case of the Ca-K line by Skumanich et aL (1984).

The group of solar-rotational peaks from November 1979 to March 1980, shown in Figure 2, involve an episode of major solar activity that was discussed previously by Donnelly et aL (1986a, b). They showed the following: (1) The episode (Figure 2) starts near the peak of an intermediate-term variation (as seen in Figure 1 in the rise of the minima in the solar rotation modulation) that started in about July 1979 and ended about March 1980. The strong amplitude of 27-day modulation from November 1979 to March 1980 comes not only from the new major active regions that produce the high rotational peaks, but also from the rapid reduction in activity at the solar longitudes near the center of the solar disk at the times of the rotational minima. (2) The first rotational peak in November was accompanied by a major dip in the total solar irradiance (Hickey and Alton, 1984) caused by new large dark sunspots (Hudson et aL, 1982). (3) F10 rises to its largest rotational peak on the first rotation and then decays to much lower peaks on subsequent rotations while the 2000 A UV flux reaches its highest rotational peak on the second rotation and decays very little on the subsequent rotational peaks. (4) The temporal evolution of the hot coronal lines, like the Fexv line shown in Figure 2, is similar to that ofF10, while the evolution of the H L~ flux~is similar to that of the 2000 UV flux. (5)The ratio of the size of the first peak (November) in the rotational modulation to that of the second rotational peak (December) progresses from less than one for the photospheric UV flux and the strong chromospheric H-L~ emission line, to slightly greater than unity for the H Lfl line, to larger values for the coronal Fexv line and F10. (6) The rotational peaks tend to be wider for the coronal emissions, like the Fexv line, than for the photospheric 2000 ~. UV flux. The day-to-day jiggles in the soft

TEMPORAL TRENDS OF SOLAR EUV AND UV FULL-DISK FLUXES

OCT NOV DEC 1979 JAN 1980 FEB 5 o o L I I ' W ' I ' q ' t ' ~ ' l l ' ~ ' I ' ~ I ' ' J5oo

~ [tt~ Daily Mean 1-8 ~, ,11 GOES|

; 4ooL"tt"\J! "'4 N 300 "i-

250 ~ / ~. ~ ~ 0~awa SGD 4 250 IE

150 L ! % ~ ,~. " 4 . ~ ]150 I 40 i- ~ !, 284 A " /~ ' . , ~ ' k ,~. 7 40

L Z " ' '~ ' % '~" "

7.5L- ,~# ', 200 ; i ~ ,,' ~ ~ -375 ~ ' ', 204 A / ~ ~ '~" ~", a q

7.o - I ,>%t / ',47.o 3.2F p.s' ' g ' ~ 'b, k,~ , ; "1 ~. ,, y "<to_l ~"

, o', \ ~ , , ,, ~ R~J65 30 [- ,178-183A ', '~,' 4 "V i I" .' , ' .... 7, �9 re FeXi ~ i ! ~ ' ; ~ W . 4 ~']~

1.8 ~ i t " - 7 ~x=,.~d ~ Fe IX " /

; I E ! . t ; , , ,

u., e ! '~Q.1216 A" # L ~ I r / ",~,~ I % % , :

'~' 2006 ,~ ' / ' I J IM ~'7 -Ol ~ t ~ #, ~ . ~ ~us-~]_Ol 0 95 ~ "~- *~/'x ~ "~ / Heath et al. (1984) n n 95 �9 J 1 2O0O- I/' ~,3 Long-term trend removed 4 "

p<~r ,!t 2050 ,~ J ~ Ca-K Based Model t .~' ~. \ ~' ~ ;,, ~ea. ota,. (t982) j

0 9 0 ~ ~l ~ / R ~ , ' ', _ , ~ o 090

150 ; ~Suun~eP?R-~ kk~. A -~ l , ,~dsGo ~ 150

lOO fl , ~ , ~ " 4 1 ~ ~ ~ , i ~, i I , [ ~1oo OCT~ NOV DEC 1979 JAN 1980 FEB

Fig. 2. An episode of major group of active regions.

41

X-ray flux are greater than in the other observations because the soft X-ray measurements are very sensitive to solar flares and other fluctuations in the hottest parts of active regions. The rotational peaks of the soft X-ray flux are much broader than for the hot coronal EUV line; the EUV rotational valley near 25 December barely exists in the X-ray flux. Note how narrow the rotational valleys are in the soft X-ray flux near 28 November and 25 January in Figure 2.


The newly added Fe IX, xI, and xII! measurements illustrate the following features: (1) they peak higher relative to the adjacent local minima on the second and third solar rotation than for the first rotational peak, (2) they tend to remain near their rotational peak values longer, before their main decline when the major active regions approach and pass over the west limb, and (3)they tend to reach their lowest values per solar rotation later and then start to rise more rapidly with the next rotational enhancement than do the chromospheric EUV and photospheric UV fluxes. These features are shown here as one example pertinent to the more general cross correlation results discussed in Section 5. The Fe IX, xI and XIII fluxes include some upper transition region emission as well as 'cool' (T< 2 x 106 K) coronal emissions. Conversely, the Fexv and xvl emissions are strictly coronal and are strongly emitted by active-region sources at T > 2 x 106 K (Jordan, 1969).

4. Persistence of Solar-Rotational Variations

Episodes of major activity rise, grow, peak, and then decay (and sometimes rejuvinate) over typically four to nine months (Harvey, 1984) and dominate the solar-rotation variations of the full-disk EUV and UV fluxes (Donnelly et al., 1986a, b). An emerging episode often involves new active regions at different solar longitudes than in the previous episode; consequently, the phase of solar-rotational variations then shifts. The autocorrelation function (AF) is very useful for studying recurrence within a time varying signal, even one that includes phase jumps between quasi-periodic episodes. AF correlates a time series with the same time series delayed or lagged by a number of days. Before computing the autocorrelation function for the solar rotational variations, the long-term trends were removed. A twelfth-order polynomial was least-squared-error fitted to each data set as a function of time during 1979-1980. Then the polynomial value for each day was subtracted from that day's observation. The residuals containing the short-term variations in their full strength are devoid of long-term variations (years) and very weak in any remaining intermediate-term variations (several months). The analyses discussed in the remainder of this section and in Sections 5 and 6 involve these same two years of residuals of short-term variations, where the numerical analyses involved are the same as those discussed by Heath et al. (1984) and Donnelly et aL (1986a, b).

Figure 3 shows AF for the chromospheric H Lflline, the hot coronal FexvI line at 335 A, and the 169-173 A_ band, which is dominated by an FeIX line. Similar results are included for the independent measures of the photospheric UV flux in the AII absorption continuum, the mountain-top observations of the He I equivalent width and F10. The strongest common feature is that peaks occur in AF near integer multiples of the solar rotation rate of about 28 days of lag. These peaks are very high in the case of the Ah UV data and are progressively lower from the top to the bottom of the figure. This indicates that the solar-rotational variations are very persistent in the UV flux and much less so in the transition region emission.

To obtain a quantitative measure of this behavior, the persistence was defined as the average of the highest autocorrelation values near each of the first four solar rotations

TEMPORAL TRENDS OF SOLAR EUV AND UV FULL-DISK FLUXES 43

Fig. 3.

"O

E <

.o_ ~a

w <

0.5 ~L_ A 205/b, AI I Gontinuum N~7~U988~ 0.5

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1026 A H Lyman /~ AE-E

o A, A . , ~ . A . . ~ . . / 1 A.,, A ,~ ~'/' ~' ~''~ " ' V 'V'v'v'v v ,.~o

-0.5 E_v ~ -0.5 1.0 -- ~ 1.0

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0.5 ~ 0.5

A, ~, A ~ , h A ,A ,A ,h o /' v v'v 'v ~ 'V ,v 'v'v v '~ ,v,~O -0.5 r - ~ -0.5

1.0 -- o ~ 1.0 i 169-173 A Fe IX AE-E : 1

0.5 -- 1979-1A~ 0.5

_ 0 5 l - I I I I I I I I I I I I " I I I I I I I I I I I 1 1 - 0 . 5 " 0 28 56 84 112140168196224252280308336364

Lag (days)

A u t o c o r r e l a t i o n func t ions for the s h o r t - t e r m va r i a t i ons in 1979-1980 .

of lag, i.e., near 28, 56, 84, and 112 days of lag. The results are listed in Table I. This persistence of solar-rotation induced variations (PSRV) is very high for photospheric UV flux variations. Using 50 A bands at UV wavelengths for this same time interval (1979-1980), PSRV was found to be very high, like that for 2050 A, throughout the 1650-2550 A_ range. In Table I, PSRV is high for chromospheric UV and EUV fluxes, medium for broadband EUV fluxes of mixed coronal, transition-region and chromospheric fluxes, low for R, F10 and hot coronal EUV fluxes, like the Fexv and xvI lines, and very low for the 1-8 A soft X-ray flux (X) and FeIx flux from the


TABLE I

Temporal and source characteristics of solar EUV, UV and solar indices

Wave- Ions Source log T Persis- Asym- Half lengths lines peak tence metry widths

(A) PSRV ACC days

335 Fexv I HC 6.4 2 4 12.5 284 F e x v HC 6.3 2 - 4 13.1 255-300 - M1 3 - 1 13.4 200-204 F e x u I CC 6.2 4 - 9 13.5 190-206 - M2 3 - 11 13.5

206-255 - M3 3 - 3 14.1 178-183 FexI T + CC 6.1 2 - 15 14.7 168-190 - M4 2 - 12 12.1 169-173 FeIX T 5.8 1 - 10 15.4 304 HeII Ch + BT 4.8 3 - 1 10.5

584 HeI Ch + BT 4.5 4 2 9.7 1026 H L/~ Ch + BT 4.3 4 - 2 10.9

590-660 M5 3 0 11.1 510-580 M6 3 3 12.2

1216 H Let Ch + BT 4.3 4 - 2 11.3

2050 AII cont Ph < 4 5 0 10.0

1-8 Soft X-rays HC >6.5 1 31 13.1 F10 10.7-cm radio Ch + T + CC + HC 2 21 11.8 R Sunspot number Ph 2 13 10.5 P Plage index Ch 3 9 - 10830 HeI Ch + BT 4 12 10.6

Source region code: HC = hot corona, strong contribution from above 2 x 10 6 K; CC = cool corona, 1 to 2 x 10 6 K; T = chromosphere-corona transition region; BT = base of the transition region, Ch = chromosphere, Ph = upper photosphere; M1 = mixture of HC and Ch + BT; M2 = mixture, T and C C ; M3 = mixture, Ch + BT, T, CC and HC; M4 = mixture, T and CC; M5 = mixture, Ch, T and CC, M6 = mixture, Ch and T.

loglo Tvalues for Fe are the temperature for the peak percentage density for that level of ionization, from Jordan (1969) and for soft X-rays from Donnelly e ta l . (1977). The emission occurs over a range of temperatures near this representative temperature.

PSRV is the persistence for solar-rotational variations measured as the average of the first four rotational peaks (lags near 27, 55, 83, and 110 days) of the autocorrelation function multiplied by ten. PSRV = 1 is very low persistence, 2 is low, 3 is medium, 4 is high, and 5 is very high indeed.

ACC is the asymmetry of the cross-correlation with respect to the photospheric 2050 A measurements measured by i00 • (CCL - CCR)/(CCL + CCR), where CCL is the max imum cross-correlation value near one solar-rotation of lag to the left in Figure 4 and CCR is that near one rotation to the right. ACC = 0 means the cross-correlation function is symmetrical. Large positive values tend to be caused by the ensemble of solar-rotational variations during episodes of major activity peaking and decaying earlier for the test data than for the reference 2050 A data and vice versa for large negative values.

Half widths are the mean width of solar-rotational variations at an intensity of half the rise to local maximum after removing a linearly interpolated background level from the preceding local minimum to the succeeding minimum.


transition region and 'cool' corona (T < 2 x 106 K). This pattern in PSRV is consistent with the pattern shown in Figure 2, where the temporal peak for F10 is high on the first rotation and much lower on subsequent rotations, which would contribute lower values to PSRV than would the more uniform series of rotational peaks in the 2050 A data. Note that the persistence peaks at Fex:II in the FeIx, xI, xIII, xv, and xw series. Both higher random instrumentation noise and higher real solar variations from the random evolution of individual active regions would also decrease PSRV in solar flux data through increasing the amplitude of random temporal variations relative to the solar rotational variations that contribute most to PSRV.

Note that the FexvI results in Figure 3 have a large increase in amplitude of the peaks every 27 or 28 days starting after the 182-day tick, i.e., after half a year. A similar effect occurs in the H L/~ data where the increase relative to the 28-day peaks that precede it is smaller. This increase in AF after 182 days is caused by the roughly in-phase alignment of the solar-rotational modulation from the peaks of one episode to another about half a year later for this particular two-year period. We conclude that this increase in amplitude is most prominent for F e x w (and Fexv, not shown) because the solar- rotational amplitudes of those hot coronal emissions in an episode of major activity tend to peak on the first rotation and decay more quickly than the chromospheric fluxes, which causes the amplitude of the autocorrelation to drop to lower values at lags of one to six solar rotations for F e x w than for the chromospheric fluxes. The peak in AF at a lag of seven rotations is similar in magnitude for both the FexvI and H L/?cases.

Puga (1984) showed that the major Ca-K plages during 1973-1983 were more persistent than either the sunspot number or sunspot area. Lean and Repoff (1987) studied five years of data (1978-1982) and showed that the photospheric 2050 ,~ flux measurements were slightly more persistent in their solar rotational variations than their 2050 A model results, based on chromospheric Ca-K plage data, and that the persistence of the chromospheric H Le measurements from the AE-E and SME satellites was intermediate between the UV observations and model results. All three of those results were much more persistent in their solar rotational variations than both F10 and the sunspot blocking factor, which causes dips in the total solar irradiance.

5. Asymmetry of Cross-Correlations

Cross-correlations of EUV fluxes with respect to the NIMBUS-7 measurements of the photospheric UV flux at 2050 ,~ are shown in Figure 4. These cross-correlations compare the short-term variations caused by the solar rotation of plages distributed inhomogeneously as a function of solar longitude. The dashed background curve in each case is the autocorrelation function for the 2050 A flux, which by definition is symmetric about the central zero-lag line. The closer the cross correlation curve is to the autocorrelation curve, the more similar the short-term EUV variations are in shape to those of the photospheric UV flux. The magnitude of the variations of course are quite different; the percent short-term variation at 2050 ,~ is typically about 5~o while that for the FexvI line at 335 A is about 200~o (max./local-min., see Figure 1). Of course


Fig. 4.

1.0 60 40 20 0 20 40 Days o f L a g

60 1.0

0.5

- 0 . 5 60 40 20 0 20 o 40 60 10,7 cm Lagged 2050 A Lagged

1.o Z"l"}"l i i 11 [ / i i i i i i ; ~ l ; i i f ~ t i [ i i i ~ ~Cross Corretation o ~, ~ --- 2050 #, Autocorrelation --Fe XVI 335 A and 2050 A ~ , -

o~ i , i i ,, ~ !", 7_

,__, , ,

60 40 ~ 20" 0 20 o 40 60 Fe XVI 335 A Lagged 2050 A Lagged

1.0 L " I I I " [ T ' C [ T T I I I I I I I ~ , I t I I I I I I I I I I I I [ I _--Cross Correlation ~ ~}~,~ ---2050 ,~ Autocorrelatir -Fe XV 284 ,~ and 2050 A ~ h

o,, ,, "A 'A 0 l i i . . . . . . . . I

-0.5 ] I J I I ]~. I I I l ~r~'[ ITI "~- ~! I t I I ! [ I I l ] JT 60 40 o 20 0 20 o 40 60

Fe XV 284 A Lagged 2050 A Lagged

0,5

-0.5 60 40 20 0 20 40 60

1-8 .~ Lagged 2050/], Lagged

t,o 111 I II t f P II I ILIf l . t . [ [ - ' [ T ~ Cross Correlation ~ ---2050 #, Autoeorret~tion

_ FeXt 178-183 A f, -- and 2050 A~' ~ ' ~

0 . 5 ; ; ; ~ ,

0 - -~ ;

-0.5 ~ 60 40 20 0 " 2~ 40 60 Fe XI 178-183 ,~ Lagged 2050 ,~, Lagged

I"0LLI I I I I I I I I I [ I I I I Is I I [ I ["iqqqTT'[-C~ I - Cer~ C69:l17t'An /'---2(150/~ A . . . . . . . . ion and 2050 ,~ ~ j ~ ~ ~ -

0,5 /i /i / , !", :

-u,~ 60 40 20 0 2b ~ 60 Fe IX 169-173 A Lagged 2050 A Lagged

1.0LLI I I I I11 I I I I I I I I I. 1171111111111111L m Cross Correlation

H ky ~ 1026 ,~

0,51 "

-i'50_, I I ] l l l l l l [ l l l l l

--- 2050 ,~, Autocorre[ation

f l l r l l [ J i i l l ~ l l 60 40 20 0 20 ~ 60 60 40 20 2b 40 60 Fe XIII Lagged 2050 A Lagged H Ly 13 1026 ,~ Lagged 2050 #, Lagged

Days of Lag

Cross-correlation functions with respect to the photospheric UV flux at 2050 A for the short-term variations in 1979-1980.

the change in energy flux (W m -2) for a 1 ,~ bandwidth is higher at 2050 ,~ because the average flux there is so much higher than in the FexvI EUV line. In Figure 4, peaks occur at zero lag and integer multiples of a solar rotation of lag, which is just a consequence of the strong solar-rotational periodicity present in both of the two data sets compared in each section of the figure.

The important features in Figure 4 are the following: (1) The cross-correlation curve


for F10 versus 2050 A and also that for X (soft X-rays) versus the UV flux are highly asymmetric to the left, with the valley near 14 days to the left being deeper than that near 14 days to the right, the peak at 27 days to the left being higher than that at 27 .days to the right, the valley near 41 days to the left is deeper than near 41 days to the right, etc. This behavior is caused by the quasi-periodic temporal structures in F10 and X tending to strongly lead that of the 2050 A flux. For example, this conclusion is consistent with the temporal behavior during the episode of major activity illustrated in Figure 2, where F10 and X tend to rise to their highest peak during the first solar rotation when the new activity dominates the full-disk flux variations and then decay quickly in the peaks on subsequent solar rotations, while the UV flux rises to its highest rotational peak on the second rotation and decays slowly in the subsequent rotational peaks. (2) The cross-correlation curve for Fexw 335 A versus 2050 A is slightly asymmetric to the left. This is interpreted in terms of the evolution of solar rotational peaks of the FexvI data tending to slightly lead that of the UV flux, which again is consistent with their behavior during the major episode of activity illustrated in Figure 2. (3) The cross- correlations of the chromospheric H L/~ emission line at 1026 A versus the UV flux are fairly symmetric and close to the 2050 A autocorrelation curve. Results for other chromospheric EUV flux measurements and also the chromospheric UV fines of Mg II H & K near 2800 A are very similar to the symmetric and highly correlated results for H LB. (4) The curve for the hot coronal Fexv emission line at 284 ]~ versus the UV flux is roughly symmetric but less highly correlated than the chromospheric fluxes. (5) The curves for Fexm, FexI, and Fe~x are asymmetric to the right, with low correlation peaks for Fe IX and xI. These low peaks are consistent with the low persistence in solar rotational variations for Fe Ix and xI. In summary, Figure 4 shows that the symmetry progresses from the amplitude being strongly asymmetric to the left for the coronal X and F10, weakly so for Fe xvI, roughly symmetric for Fexv, and asymmetric to the right for Fe xIII, xI, and Ix, then switching to a much higher correlated and symmetric result for the chromospheric EUV and UV fluxes. Table I quantitatively summarizes the amplitude asymmetry. Note that the order of increasing absolute magnitude of the asymmetry factor in Table I quantitatively summarizes the amplitude asymmetry. Note that the order of increasing absolute magnitude of the asymmetry factor in Table I is roughly similar to that of the increasing magnitude of the persistence factor from the autocorrelation function.

The widths of the cross-correlation peaks in days of lag shown in Figure 4 exhibit a progression from being several days of lag wider than the UV autocorrelation curve for X, a few days wider for F10, and FexvI, xv, xIII, xI, and Ix, and nearly the same for the chromospheric EUV flux ofH Lfl. This is interpreted as being caused by the 27-day solar rotational variations tending to decrease in the temporal widths of their peaks (increase in the widths of their valleys) from the X-ray flux, to the Fe lines and to the chromospheric and photospheric fluxes. A study of the half-widths of solar rotational variations for these same data (Donnelly et al., 1986b), which is summarized in Table I, gave a similar order, except that the average half widths for Fe rx and xI were slightly higher (but not statistically significantly higher) than the value for the solar soft X-rays.


Considering the strong short-term temporal variation of the solar soft X-ray flux caused by the temporal evolution of active-region hot spots (and not due to solar rotation variations), which are apparent in Figure 2, the half-width study may underestimate the half width of the soft X-ray variations caused by solar rotation.

6. Episodic Piage Deficiency

Figure 5 shows curves of amplitude as a function of time from complex demodulation analysis of the short-term variations (Chapter 7, Heath etaL, 1984). This analysis evaluates the amplitude and phase in the local vicinity of time t of the sinusoidal signal within each data set at a selected frequency or period. The period chosen was the 27.7-day solar-rotational peak found in the power spectra for these 1979-1980 EUV data (Donnelly et aL, 1986a). A peak in one of these curves means that the solar-rotation modulation in that data is stronger within several weeks about that time than at other

E <

re c- O

"0 0 E q X

E 0 0

Fig. 5 .

2( N

I:

1979 1980

Complex -demodu la t i on ampl i tude for a 27.7-day per iodic i ty for 1979-1980.


times. The two main peaks for the photospheric 2050 A data in December 1979 and July 1980, have later, broader and slower decaying peaks with fewer subpeaks than for the coronal F10 and FexvI data and chromospheric P. The peaks for June and December 1979 for the HeI line at 584 ~l are even broader than for the 2050 A data. Many of the trends as a function of time and wavelength in Figure 5 are just another view of the same trends previously discussed for episodes of major activity (Figure 2), persistence of solar rotational variations (Figure 3), and asymmetry in cross-correlations (Figure 4). The main new point that Figure 5 emphasizes is that the solar-rotation modulation of the Ca-K plage index P peaks earlier and decays sooner than the chromospheric EUV and photospheric UV fluxes. Similar results for P relative to the measurements of the chromospheric HeI line at 10 830 it and NIMBUS-7 observations of photospheric UV fluxes occurred for the major episodes of activity of October 1981 through March 1982 and June-October, 1982 (Puga et aL, 1986). This result is consistent with (a) the persistence for P in Table I being lower than for most chromospheric EUV fluxes and photospheric UV fluxes and (b) the asymmetry function being intermediate between that of FexvI and F10 and much higher than that of the photospheric UV flux and most chromospheric EUV fluxes.

Another systematic progression in the short-term EUV temporal variations shown by Donnelly et al. (1986a, b) is that the ratio of the power in lines near 13 days and 27 days in the power spectra of the 1979-1980 short-term temporal variations is high (�89 for the photospheric UV flux, medium (-~) for the chromospheric EUV flux, and small to negligible for the hot coronal EUV fluxes. This progression is caused by the 13-day periodicity being very sensitive to the way active region emission varies with the solar central angle as the region rotates across the solar disk. Analysis and modeling of this behavior in full-disk UV and EUV fluxes has been completed and will be published later in a separate paper. Solar measurements with spatial resolution have not yet been used to study the average center-to-limb dependence of active region emission at EUV and UV wavelengths.

7. Ratio of Long-Term to Short-Term Variations

The average valley-to-peak amplitude of the short-term variations appears to vary relative to the long-term variations from one wavelength to another in Figure 1. To obtain a quantitative measure of this property, the amplitude of the long-term enhancement was divided by that of the average short-term variation and the results are shown in Figure 6. The long-term peak flux was measured by the 13-month mean of monthly means (half weight for months 1 and 13), for February 1979. That was the maximum such value for most of the fourteen AE-E wavelength groups studied. The monthly mean at the start of the data studied, July 1977, was subtracted from the long-term peak flux to provide one common measure of the long-term increase for solar cycle 21. The amplitude of the short-term variations was measured by the root mean square (r.m.s.) deviation of the 1979-1980 short-term variation, after the twelth order polynomial was used to remove the long-term trends.

50 R I C H A R D F . D O N N E L L Y

11 O3 e-

.o 10 t~

r

g8 o x: 7

o ~ 6 r -

.s ~ 5

r

~ 3 C

--q 2

o % m

0

Fig. 6.

Short-Term Var ia t ions M e a s u r e d by R M S R e s i d u a l s

I I I J ] I I ] I I I I I ] I o l ] ] ]

X

X o

o

x

[] �9

o +

[]

[]

�9 AE-E Instr. No. 1 O AE-E Instr. No. 2 x AE-E Instr. No. 5 E l - + AE-E Instr. No. 9

- - A NIMBUS-7 SBUV D Ground-Based M e a s u r e m e n t s

I I I I I I I I f I I I E I I I I ] I -

d~ cb eJ ~ c~

e4 O

Ratio of long-term variations to the r.m.s, short-term variations.

Figure 6 shows that the amplitude ratio of the long-term to short-term variations vary systematically with wavelength. The order of the entrees in Figure 6 is somewhat arbitrary. The first wavelength, 2050 A, is emitted from the upper photosphere, which is from a deeper solar altitude than all the other full-disk fluxes shown. This one result involves an estimate of the long-term variation at 2050 A based on the correlation between four years of 2050 ~ and 10 830 ,~ flux determined by Donnelly et al. (1985). The H Lfi line and the He I and n lines are emitted from the chromosphere and base of the transition region, and the sequence of Fe fines are from the middle transition region to the hot coronal lines of Fexv and xvI. After the emission lines and narrow bands dominated by a line (like FeIx, xI, and xm) are the broadband measurements, which involve a mixture of chromospheric, transition region, and coronal emissions. These bands are ordered by decreasing wavelength from left to right because the mixture of emissions was thought to range from predominantly chromospheric to predominantly transition region and coronal emissions in that order.

The important results in Figure 6 are the following occurrences: (1) high ratios for the


Fe xIII and Fe xI lines from the upper transition region and cooler portions of the corona (T < 2 • 106 K), (2) moderate ratios for the chromospheric fluxes of He n at 304 A, He I at 584 A, and 10 830 A, and H L/~ at 1026 A, and (3) low ratios for F 10, sunspot number R, Ca-K plage index P, and the UV flux at 2050 A. The broadbands 168-190 and 190-206 A include the FexI and Fexm results, respectively, so it is not surprising that their ratios are also high. Note the close agreement between the AE-E measurements of the chromospheric fines at 304, 584, and 1026 ]~, and the independent ground-based measurements of the 10 830 .~ line. Observe also that the photospheric 2050 ,~ ratio and the coronal FexvI and F10 ratios are lower than the chromospheric results, R has an even lower ratio, and the chromospheric P has the lowest value of all. The ratios for wavelength bands that include a mixture of chromospheric, transition-region and coronal emissions have the high ratios. The FeIx, xI, xIII, xv, and xvI ratios form a sequence that peaks with FexIII.

8. Analysis Summary

Several progressions in the temporal characteristics of full-disk solar UV and EUV fluxes have been identified as a function of the solar source region, including the following:

(1) Episodes of Major Activity. During episodes of major activity, the 27-day periodicity of the chromospheric 10830/~ He I line, the 2050 A photospheric UV flux and the chromospheric L~ emission line tend to reach their maxima on a later solar rotation and then decay more slowly than for the cases of the solar 10.7 cm flux (F10), the sunspot number (R) and the Ca-K plage index (P). The EUV temporal variations progress from the chromospheric H L~ being similar to that of the photospheric UV flux, through H Lfl, HeII 304 A, to the coronal Fexv and xvI lines varying more like F10. See Figures 2 and 5.

(2) Persistence of Solar-Rotation Variations. The persistence of solar-rotation induced variations is as follows: (a) very high for photospheric UV flux variations, (b) high for chromospheric UV and EUV fluxes, (c) medium for P and broadband EUV fluxes of mixed coronal, transition-region, and chromospheric fluxes, (d)low for R, F10 and coronal EUV fluxes, like the Fexv and xvi lines, and (e) very low for the hot coronal 1-8 A soft X-ray flux (X) and the FeIx flux from the transition region and 'cool' corona. See Figures 3 and 5 and Table I.

(3) Cross-Correlation Asymmetry. The cross-correlations of short-term variations (days, weeks) with respect to the photospheric 2050A flux progress as follows: (a) strongly asymmetric in an episodically leading direction for X and F10, (b) moderately asymmetric for R, P, and the 10 830 A HeI line, (c) quite symmetric for chromospheric EUV and UV fluxes and photospheric UV fluxes, and (d) moderately asymmetric in a lagging direction for transition region and 'cool' coronal fluxes from emission lines of FeIx, xI, and xIII. See Figure 4 and Table I.

(4) Half Widths. The average temporal width at half the rise from local minimum to maximum during solar-rotation induced variations increases from about ten days for the


photospheric 2050 ~ flux, to about eleven days for chromospheric EUV fluxes, to thirteen days for coronal Fexv and xvI lines and the soft X-ray flux, to fourteen to fifteen days for FexI and IX lines. See Table I.

(5) Plage Deficiency. The plage index P is deficient in its temporal variations relative to most of the chromospheric full-disk fluxes in the following ways: a) lower persistence in solar-rotational variations, (b) asymmetrical cross-correlations with respect to the UV flux rather than high symmetry, (c) earlier peak and decay in the complex-demodulation amplitude for a periodicity of 27.7 days, and (d) much lower ratio of long-term to short-term variations. See Table I and Figures 5 and 6.

(6) Long-Term~Short-Term Ratios. Ratios of the long-term amplitude increases during the rise of solar cycle 21 to the amplitude of short-term solar-rotation variations are higher for the AE-E measurements of full-disk EUV fluxes that originate in the chromosphere and for ground-based measurements of the chromospheric HeI absorption line at 10 830 A than for F10, R or P. Lower ratios occur for the photospheric 2050 measurements and hot coronal FexvI flux at 335 A than for the chromospheric fluxes. The highest ratios occur for the transition-region and coronal lines from temperatures below about 2 x 106 K, like the FexIn line at 202 A. See Figures 1 and 6.

9. Discussion

The above trends will be discussed in terms of the temporal evolution of major groups of active regions and the third component, i.e., the collective influence of numerous sub-plage enhancements. Some of the explanations suggested below are plausibility arguments or speculation and have not been proven. Research of solar physics measurements with moderate or high spatial resolution, rather than full-disk measurements, is the proper avenue to prove some of these suppositions. Questions are identified for further research.

9.1. EVOLUTION OF MAJOR ACTIVE-REGION GROUPS

The results of Figures 2-5 are related to the same physical cause, the evolution of emission as a function of the. source-region temperature for major groups of active regions. During an episode of major activity, the chromospheric EUV and photospheric UV fluxes rise and decay more slowly in Figure 2, which causes high persistence in Figure 3 and symmetric results in Figure 4. These results for chromospheric and photospheric enhancements are consistent with old ground-based results that showed chromospheric plage and the associated photospheric faculae were quite long lasting relative to the lifetime of the associated sunspots (Kiepenheuer, 1953; Puga, 1984). The hot coronal emissions rise rapidly to a peak and quickly decay over several solar rotations, which contributes to low persistence in the solar-rotation peaks of the autocorrelation in Figure 3 and the asymmetry in Figure 4. These coronal Fe line and radio results in Figure 2 are consistent with the early results of Neupert (1967). The FelX and xI fluxes rise slower than for Fexv and xvI from one rotational peak to the next in Figure 2 but decay quicker than the photospheric UV flux and the chromospheric EUV


fluxes and include more temporal variations randomly related to the solar-rotation induced variations causing low persistence in Figure 3 and asymmetry to the right in Figure 4.

The new features in Figures 2-5 can be expressed by the following questions: why is F10 more asymmetric in an episodically leading direction in Figure 4 than either F e x w or xv? Why are the FeIx-xIn results asymmetric in an episodically lagging direction relative to the chromospheric EUV and photospheric UV fluxes and why do they contain more short-term variations randomly related to the solar rotation variations than either the Fexv and xvI hot coronal emissions or the chromospheric EUV fluxes? Why is the persistence for the Ca-K plage index P so low or the asymmetry so high in Table I relative to the other chromospheric measurements?

9.1.1. Centimeter Radio Emission

During the past decade, solar radio astronomy at centimeter wavelengths has concen- trated on high spatial resolution measurements. Small bright emission cores have been observed and interpreted as resulting from gyroresonance absorption and emission (Alissandrakis et al., 1980). Such observations have been used to estimate magnetic field strengths in the active region corona (Kundu et aL, 1980). Some bright radio emission regions have been suggested to be caused by nonthermai processes (Webb et al., 1983). The large diffuse radio sources associated with active regions are interpreted as thermal bremsstrahlung emission (Felli et al., 1981). Meanwhile, the full-disk 10.7 cm flux (F10) is the most frequently used solar radio measurement, i.e., used by terrestrial scientists to roughly estimate the effects of temporal variations in the solar EUV and UV fluxes, with no consideration of the recent solar research on bright gyroresonant emission cores.

Recent studies of high spatial-resolution centimeter data do not compare the net radio emission from these small bright cores with the total active region emission, where the latter includes the large diffuse thermal bremsstrahlung sources. Felli et al. (1981) found that the bright core emission at 6 cm for one active region was less than 5 ~o of the total flux from that region; but that was an old third-rotation region, one that contributed partially to the rise of the small solar rotational variation in October 1979 in Figure 1. The frequency of flares is high during the rapid growth phase of active regions. We speculate that the gyroresonance and possibly nonthermal radio emissions are stronger relative to the total region's emission during the first solar rotation of an episode of major activity, i.e., at the peak of the coronal emissions like that of the F e x w line, than later during the decay of the episode. This may help explain the stronger asymmetry of F10 with respect to the 2050 A photospheric flux than for the FexvI flux in Figure 4. Questions for further research include the following: What is the distribution of the ratio of net gyroresonance emission and/or nonthermal emission) at centimeter wavelengths for active regions as a function of time during the evolution of an active region or collectively during an episode of major activity? How do these results vary with wavelength? What are the consequences at 10.7 cm? Considering that the quiet-Sun magnetic field is now known to consist of small bundles of strong fields, what are the conse-


quences on the classical interpretation of the quiet-Sun flux at centimeter wavelengths, which neglected the effects of magnetic fields, and the consequences for active network emission at centimeter wavelengths?

9.1.2. Cool Coronal and Transition Region Emission

A simple explanation of the half-width for solar rotational variations of the FeIx, xI, and xtII lines being broader than those for the Fexv and xvI lines may be that the longitudinal extent of the "cool' (T < 2 x 106 K) coronal and upper transition region emissions of active regions may be greater than that for the hot coronal emissions of major groups of active regions. At transition region wavelengths, bright solar plumes are observed overlying sunspot umbrae (Noyes et aL, 1985). On the other hand, Schrijver etaL (1985) found that 'areas of active regions determined from chromospheric emissions are systematically larger than areas determined from transition region emissions'. They explained this as being caused by the magnetic flux tubes central to a plage having much thicker transition regions than those near the edge of a chromospheric plage. Transition-region and cool coronal emissions may reach their maximum later in an episode of major activity both through the cooling of formerly hot corona loops (Neupert, 1967) and through the formation of new small regions among the magnetic fields of former active regions (Garcia De La Rosa, 1984). The combination of random walk and differential rotation may spread old emission regions more than for the short-lived hot coronal sources. Cool coronal loops are probably more involved in inter-region connections. The combination of lead-spot plumes, plage related sources and old region loops is consistent with a greater longitudinal extent for the FeIx, xI, and xIII emissions than for chromospheric or hot coronal emissions. Also, the dependence of the plage emission for these FeIx, xI, and xgI emissions on the solar central angle may be broader through less absorption in the solar atmosphere since these wavelengths are shorter than for the Fexv and xvI lines. Do EUV measurements with spatial resolution substantiate this? Do the regions in the trailing end of the group tend to be older than those at the leading edge? Do the region-to-region coronal connecting loops contribute to this effect?

Neupert's (1967, Figure 6) early results of a fast evolution for an Fexvt line, somewhat slower evolution for Fexv, and a much slower one for FexIH are consistent with the high persistence for FexIII and low persistence for FexvI and xv in Table I and the shift from asymmetry to the left for FexvI to asymmetry to the right in Figure 4. Neupert's results suggested that hot coronal active region loops may evolve into or produce cool coronal loops. The above discussion indicates that quite different spatial regions are involved in the evolution of full disk fluxes. The slow evolution of cool coronal emissions of Figure 2 and Neupert (1967), like the long-lived chromospheric plages, would suggest the persistence of solar-rotational variations (PSRV) should be high for these cool coronal emissions. This is constistent with the high persistence for FexIII; however, Table I shows the persistence is very low for FeIx and low for FexI. Fe Ix is the most affected and Fe xIII is the least affected of these three lines by transition region emission. Part of the low persitence may result from higher observational noise


for Fe IX. However, we suggest that much of the cause of the very low persistence for FeIx and low persistence for FexI is that the transition region portions of these emissions have low persistence, which is partly due to temporal variations in the emission measure in the 105 to 10 6 K range that is rather independent of the full-disk variations at higher and lower temperatures (Chapter 3, Donnelly et aL, 1986b; Figure 2, Neupert, 1967).

9.1.3. Plage Remnants

The most interesting aspect of the Ca-K plage index P is that its temporal variations differ so much from the observed chromospheric full-disk fluxes (EUV, UV, and 10 830 A) and photospheric UV fluxes. The persistence for P is low and the asymmetry factor is high in Table I, which are consistent with the main peaks in Figure 5 peaking and decaying earlier than for the next five curves above that for P. Systematic biases in the plage data might contribute to this problem; the use of an estimate of the intensity of the brightest part of the plage rather than a measured average intensity might tend to overestimate the plage importance during its rising phase and underestimate the plage importance late in an episode of major activity. The processing of plage data changed from McMath-Hulbert Observatory to Mt. Wilson Observatory on 1 October, 1979, which may have introduced some changes. We suggest that one main cause of the temporal difference of P in Figure 5 from the curves for chromospheric fluxes is the natural tendency to not scale very small features, e.g., plagettes, plage remnants, and the active network. Garcia De La Rosa (1983, 1984) found that small active regions tend to emerge in the remnant fields from previous large active centers, so the formation of new small regions may occur preferentially late in an episode of activity. The temporal differences for P in Figure 5 then suggests that the net importance of these small but numerous features is greatest late in an episode of major activity and that they have not yet spread sufficiently in solar longitude to remove their contribution to the solar rotational modulation.

9.2. THE THIRD COMPONENT

Lean et al. (1982) invoked a third componend, in addition to the well-known quiet Sun and plage components, to explain the temporal variations of the solar ultraviolet flux. The short-term solar-rotational variations were approximately explained by the second component of the model based on Ca-K plage observations; however, that component was markedly insufficient to explain the observed long-term variations, which is analogous to the ratio for P in Figure 6 being much lower than that for the UV flux. So the third component was needed to explain the observed additional solar-cycle enhancement. Skumanich et al. (1984) found the third component was required to explain the full-disk Ca-K flux over solar cycle 21. Lean et al. (1982) sometimes referred to the third component as the 'active network' and solar-rotational variations for the third component were neglected. The third component probably involves a combination of several solar phenomena called by the following overlapping labels: active network, ephemeral regions, plagettes (too small to scale as plages), plage remnants, etc. In Section 9.1.3,


we suggested that the distribution with respect to solar longitude for plagettes and early plage remnants is not always uniform, i.e., their contribution to solar-rotation modulation is important late in an episode of major activity.

Consider the results in Figure 6 and their possible link to the third component. A higher than average ratio for a particular wavelength might be achieved by the following: (1) increasing the long-term variation beyond that caused by active regions (the regions scaled in Ca-K line plage data) by increasing the third component, (2) decreasing the amplitude of the short-term variations by widening the central-meridian-distance (CMD) dependence of the active region emission, and (3) a combination of the above two processes. Wide CMD dependences like that of F10 cause weak 13-day variations and weak harmonics of the 27-day variations, but they also cause strong 27-day modulations. While F10 has a broad CMD dependence (Donnelly et aL, 1986a), it does not have a high ratio in Figure 6. So a CMD dependence as wide as that of F10 is not enough to contribute much to having a high long-term to short-term ratio. Since the ratios for chromospheric fluxes are higher than for F10, yet the chromospheric fluxes have narrower CMD dependences than FI0 (Donnelly et al., 1986a), then the first case is our main candidate for explaining the higher chromospheric ratios, e.g., through active region remnants and active network (the third component) that are spread sufficiently uniformly in longitude that they do not contribute much to the short-term solar- rotational modulation.

We speculate that the very high ratio for the cool coronal Fexm line is best explained by the third case, i.e., a combination of large third component and a very broad CMD dependence. Since the transition region and corona lie above the chromosphere and are much less dense, their CMD dependence should be much broader than for chromospheric emissions. In Section 9.1.2, we also suggested that the longitudinal extent of Fexm sources may be greater for groups of major active regions than for the hot coronal emissions. These effects combined would cause very broad solar rotational variations, which is consistent with the large half-width result for Fe xH1 in Table I, and thereby reduce the average amplitude of solar rotational variations. Spatial EUV measurements, like those of ATM Skylab, should be studied to determine whether they support or dispute these specualtions.

10. Conclusions

The progressions in temporal variations of the full-disk EUV and UV flux, summarized in Section 8, suggest large-scale organization in groups of active regions, their remnants and other small but numerous emission features. The remnants of plages and other sub-plage brightenings have been suggested to contribute significantly to the solar rotational variations late in an episode of solar activity and to be rich in chromospheric EUV and cool coronal emission, like that from Fe xm, but weak in hot coronal emission, like the Fe• and xvI EUV lines or soft X-rays. Numerous small brightenings, the so-called third component, have been invoked to explain the large long-term increases of chromospheric, transition-region, and especially the cool coronal emission of Fe xm.


It is now time that spatially resolved solar measurements be used to substantiate or disprove, this third component. It is also clear that full-disk EUV measurements should be continued as well as spatially resolved measurements until there is close agreement between the two types of observations and a better understanding of their long-term variations.

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temporal trends of solar euv and uv full-disk fluxes

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