F e x l v AND O T H E R L I N E I D E N T I F I C A T I O N S IN T H E S O L A R
E U V S P E C T R U M
S. O. K A S T N E R
Mathematical Science Consultants, hlc.. Greenbelt, MD 20770, U.S.A.
(Received 27 December, 1990; in revised form 6 May, 1991)
Abstract. A number of solar extreme ultraviolet lines, previously unidentified or given assignments at the time of observation, are found to be coincident in wavelength with transitions recently identified as aluminum-like or magnesium-like in the laboratory. Three tables summarize the results of this comparison, which does not imply positive identifications but suggests that several assignments may be valid. Additional assignments are proposed in other ions and a revised table of predicted and observed forbidden line wavelengths in the Ne ~ sequence is included.
1. Introduction
Many observed lines in the extreme ultraviolet solar spectrum have remained unidenti- fied or as tentative identifications because of a lack of laboratory-confirmed assign- ments, especially in the heavier isoelectronic sequences. This is the case, for example, in the aluminum-like spectra which involve three-electron configurations such as 3s3p3d. Recently Redfors and Litzen (1989; abbreviated as RL) were able to determine most of the levels in this configuration, except for those of the 4F term. Attention is drawn in this communication to the fact that the wavelengths of a number of unidentified lines in quiet Sun and flare spectra are comparable with the wavelengths obtained by RL and in other recent laboratory work. The results of this comparison are given below in three tables. In many cases where previous identifications have been proposed the present assignments may supplement those identifications, in the sense that both com- ponents cannot be ruled out at this time as contributing to the observed line. In certain cases it appears that the present assignment is more probable on the basis of elemental abundances or other criteria mentioned below, e.g., when iron lines are proposed to replace aluminum or argon lines. Secondary aims are to clarify previous work on the aluminum- and neon-like ions, and to note some observational anomalies.
2. Background
A review of early developments in analysis of observations of the solar EUV spectrum beow 400 A. was given by Widing (1973) who described problems involved in arriving at some of the less obvious identifications. Difficulties were generally caused by lack of laboratory data or by lack of reliable theoretical predictions, two principal obstacles which persist in some measure today.
The first prediction of solar line intensities to be expected from FexIv was that of Blaha (1971 ; MB) who included five configurations in a 47-level calculation. This work,
Solar Physics 135:343-351, 1991. �9 1991 Kluwer Academic Publishers. Printed in Belgium.
344 s . o . KASqNER
though based on early atomic data, appears still to be the only comprehensive calcula- tion of allowed line intensities in this ion. A calculation of line intensities for the Al-like ions by Kastner (1981), using the data of Farrag, Luc-Koenig, and Sinzelle (1980), treated the stronger resonance lines. More recently the results of additional theoretical calculations of atomic data for this sequence have been published by Huang (1986) and Froese Fischer and Liu (1986; FL). Observationally, however, only levels of the lower configurations have been determined, the higher configuration 3s3p3d and its levels remaining mostly undetermined until the recent work of RL.
The stronger lines of Fe xIv and NixvI are present in the quiet Sun spectrum of Behring et al. (1976; BCFD) and in the flare spectrum of Dere (1978; KD), implying that certain of the weaker unidentified lines, at least in the flare spectrum, may also be due to these ions and others in the AI Mike and Mg Mike sequences. The solar spectra were, therefore, compared with the RL list, and with other likely transitions chosen with the aid of the tabulations of Huang and FL. Recent laboratory analyses of the Mgl sequence by Churilov et al. (1985) and Litzen and Redfors (1987) also are used for comparison. A number of new assignments can be proposed as indicated in the tables.
3. Tables
For reference the energy levels of the A11 sequence are listed in Table I with the Blaha numbering. To give an idea of the kind of wavelength agreement existing between solar and laboratory observations, Table II compares the BCFD and RL wavelengths for identified FexIv lines, showing that differences generally increase as the lines become weaker, as one would expect. The difference of 37 mA for the line 289.1 A is particularly large. Such possible differences for lines which are still weaker may be kept in mind in assessing validity, i.e., a difference of 40 mA, - 0.04 A between a solar wavelength and a laboratory wavelength does not necessarily mean that the wavelengths are not due to the same transition, though a positive identification cannot be made. However, most of the wavelengths listed for unwidened lines in the tables agree to within a considerably smaller A2 -~ + 20 mA.
Table III lists significant quiet-Sun matches found in the AII sequence. In this and the following tables there are several classes of wavelength agreement: (a)close agree- ment (IA21 < 20 re?t) in which the solar line is weak, narrow and unidentified, or matches in which the solar line is wide and unidentified; (b) matches in which a wide solar line has been given a previous identification; (c) matches of close wavelength agreement in which a weak, narrow solar line has been given a previous identification; and (d) matches between unidentified solar lines of appreciable intensity and lines of elements of lower abundance (Cr, Mn). The probability of correct assignment decreases approximately in this order. Table IV lists quiet-Sun identifications in the Mg I sequence, with one assignment to a pair of lines in OII. These fall generally into the four classes of reliability, except for the line footnoted 'e' which is mentioned below.
Wavelength matches of A1 Mike and Mg Mike transitions with flare lines in the list of KD are listed in Table V. In this case wavelength agreement is less decisive because
Fexlv AND OTHER LINE IDENTIFICATIONS IN THE SOLAR EUV SPECTRUM
TABLE I
Aluminum-like energy' levels a b
345
Level Configuration Term J Level Configuration Term J
1 3s23p 2p 1/2 21 3s3p3d 4F 9/2 2 3/2 22 4D 1/2 3 3s3p 2 4p 1/2 23 3/2
4 3/2 24 5/2 5 5/2 25 7/2 6 2D 3/2 26 4p 1,,'2
7 5/2 27 3/2 8 2S 1/2 28 5/2 9 2p 1/2 29 2F 5/2
10 3/2 30 7/2 11 3s23d 2D 3/2 31 2F' 5/2 12 5/2 32 7/2 13 3p 3 20 3/2 33 2D 3/2
14 5/2 34 5/2 15 4S 3/2 35 2D' 3/2 16 2p 1/2 36 5/2 17 3/2 37 2p 1/2 18 3s3p3d 4F 3/2 38 3/2 19 5/2 39 zp, 1/2 20 7/2 40 3/2
a Numbered after Blaha (1971). b Primed terms of the 3s3p3d configuration have ~P parentage.
TABLE 1I
Comparison of BCFD and RL wavelengths tbr stronger Fex lv lines
Transition ~ MB BCFD RL j A2r i - j line 2 (/k) 2 (,'Tk) m ~
1-I1 1 211.316 211.331 15 1-8 2 274.203 274.203 0 2-10 3 264.787 264.785 2 l -6 4 334.171 b 334.171 0 2 - l l 5 220.082 220.076 6 2-9 6 270.524 270.511 13 1-9 7 257.392 257.377 15 1-10 8 252.197 252.188 9 6-31 9 c _ _
2-12 10 219.123 219.136 13 2-7 13 353.833 b 353.829 4 2-8 14 289.160 u 289.123 37
Keyed to Table I. u Identifications by Widing, Sandlin, and Cowan (1971).
Not observed; see text.
346 S. O. KASTNER
TABLE III
Aluminum-like line identifications
Ion Transition a MB RL BCFD i - j line 2 (~)
2 (~) Ident. Comment
Fexrv 7-31 17 192.638 192.630 - weak 7-32 25 193.746 193.715 - weak 7-38 21 197.776 197.847 - wide 5-25 - 216.928 216.90 SivIIt very wide 7-30 - 218.169 218.179 SxIl wide
10-36 18 223.220 223.202 - wide
5-28 - 223.261 6-29 - 224.354 224.346 - medium 8-38 - 226.040 226.017 - weak 7-17 - 290.747 290.710 - blend 7-14 - 358.681 358.67 Fexl weak
Nixvl 2-10 - 232.483 232.58 Hell wide 11-35 - 235.776 235.79 - weak 11-37 - 239.055 239.03 - weak 1-8 - 239.510 239.52 - weak 7-14 - 313.724 313.734 Mgvm
Mn xIII 7-32 - 207.130 207.112 - wide 4-27 - 227.445 227.479 - weak 9-39 - 237.369 237.333 He H wide 7-34 - 257.234 b 257.262 - strong 1-10 - 272.154 272.15 - weak
Crxn 6-31 - 220.890 220.870 - weak 5-24 - 249.374 249.388 -
10-36 - 256.370 b 256.38 Six
Keyed to Table I. b Blend.
the K D wave l eng ths have m o r e uncer ta in ty , A2 "- + 30 mA. T h e s e c o m p a r i s o n s wi th
flare lines are o f in teres t , however , s ince the efficiency o f exc i ta t ion to h igher levels,
w h i c h d e p e n d s on t e m p e r a t u r e and densi ty , is one cr i te r ion for the e s t ima t ion o f
p robabi l i ty o f co r r ec t ident i f icat ion.
O t h e r cr i ter ia w h i c h were t a k e n in to accoun t , in a s se s s ing poss ib l e a s s i g n m e n t s , were
the degree o f ion iza t ion , i.e., w h e t h e r lines o f ions o f s imilar ion iza t ion po ten t i a l s were
p r e sen t or no t in the line l ists; the p r e s e n c e or a b s e n c e o f p o p u l a t e d g r o u n d levels in
t r ans i t ions , the fo rmer cond i t i on add ing weight to a given a s s i g n m e n t ; a n d a b u n d a n c e
c o n s i d e r a t i o n s as n o t e d above in ca tegory (d). T h e s e cri teria, he re app l ied quali tat ively,
can be appl ied m o r e quant i ta t ive ly w h e n re levan t ca lcu la t ions o f e x p e c t e d level p o p u -
la t ions in F e x w a n d o the r ions are avai lable (see below). O n the o b s e r v a t i o n a l side,
t e m p o r a l co r re l a t ions b e t w e e n in tens i t ies o f l ines be long ing to t he s a m e ion, e.g.,
b e t w e e n the w eake r Fe XlV lines p r o p o s e d here a n d the s t ronger F e x I v r e s o n a n c e l ines,
c a n be conc lus ive w h e n p r e sen t ( examples are the solar analys is o f F e l d m a n , D o s c h e k ,
Fexlv AND OTHER LINE IDENTIFICATIONS IN I-HE SOLAR EUV SPECTRUM
TABLE IV
Additional line assignments
347
Ion Transition 2 (A) BCFD
(A) Ident. Comment
Fe xv 3p2( tD 2) - 3 p 3 d ( 3 D 3 ) 229.7442 229.748 3p2(ip~ ) _ 3p3d(3D2) 231.47 u 231.444 3pZ(3po) - 3p3d(3Di ) 233.46 b 233.445 3p2(lD2) - 3p3d(~D2) 257.127 a 257.136 3s3p(3p2) - 3p~(3p2) 304.894 a 304.853 d 3s3d(~D2) - 3p3d(~P~) 319.70 b 319.830 3s3d(~D2)- 3p3d(tF3) 332.854 ~ 332.77 3s3p(Ipl) - 3p2(3P2) 434.98 b 434.932
NixvIl 3p2(lD2) - 3p3d(3D3) 202.046 a 202.044 3s3p(3P~ ) - 3p2(3p2) 251.949 ~ 251.953 3s3p('el) - 3p~(Iso) 285.619 ~ 285.600
C a Ix 3s3p(3p2) - 3 p Z ( 3 P a ) 515.56 b 515.57
O 1I 2p3(2D5, ,2) - 2p~(2D5 2) 718.504 ~ 718.53
2 p 3 ( 2 D 3 , 2 ) - 2ff4(2D3:2) 718.566 ~
- weak H e II weak - weak S x weak - weak Si VIII wide
AIx MgvIt
Fe xII1 strong Fe XlXI strong S xI weak
He I
weak
a Churilov et al. (1985). b kitzen and Redfors (1987). ~ Eriksson (19877. d This line is a blend with Mn xlv. e See text.
and Seely (1985) and the labora tory analysis of Kastner , Behring, and Cohen (1975));
and in the present case may be verifiable in original N R L quiet and flaring solar spectra,
as noted by a referee. The EUV assignments suggested here rely primari ly however on
the high degree of wavelength agreement between many of the newly-identified labora-
tory lines and the unidentified, or tentatively identified, solar lines in the B C F D and K D
lists.
At longer wavelengths, predict ions of forbidden line wavelengths in the Ne I sequence
associa ted with the transi t ion 3P o --, 3p~ within the 2p53s configuration, 3P o being a
metas table level, were made by Kas tner (1982, 1983; abbreviated as S K). These were
based on isoelectronic interpolat ion from a tentative coronal identification of the tran-
sition in F e x v n . Later work by Feldman, Doschek, and Seely (1985) and Buchet et al.
(1987) has shown that the values in Table 4 of S K should be replaced by the improved
values listed in Table VI. (For F e x w I the observed coronal wavelength, 1153.20 ,~,,
places the 2p53s(~Po) level at 5951500 c m - l, if the value of 5864770 e r a - x (Buchet
e t a l . , 1987) is accepted for the level ~PI, rather than at the value o f about
5950900 c m - l.) A weak coronal line at 1749.6 ~, was recorded above the solar limb by
Fe ldman and Doschek (1977) which may be the C r x v line expected near 1750 A. This
is included in Table VI only as a tentative identification since, despi te the relatively high
348 S. O. KASTNER
T A B L E V
Line identif icat ions in solar flare spec t rum of KD
Ion Trans i t ion 2 (A) ~ K D flare
2 ( k ) Ident .
Ni x v l i
Ni XVII
Fe X[V
Fe x tv
C r x I l
Fe XlV
Fe •
Fe XlV
Fe x v
Fe x tv
F e x I v
M n x l l l
Fe x v
Fe x v
Fe x v
Ni x w
MnXl l t
Ni x v l
F e x v
Ni xvI
Fe x v
C r x l l
Ni xvH
Mn XHI
C r x I I
Ni xvI
Fe x v
Fe x v
Fe x i v
C a v I l l
C a v t l l
T i x
Cav l I I
C a w n
Ti •
T i x
Ca Ix
3f f2 (1D2) - 3p3d(3D3) 202.046 b
3s3p(1pl) - 3s3d(tD2) 215.905 r
3s3p2(4Ps,,2)- 3s3p3d(4Dv.,2) 216.928
3s3p2(2D5 2 ) - 3s3p3d(2F7,2) 218.169
3s3p2( 2D3:2) - 3s3p3d( 2F:,,2) 220.890
3s3p'?-(2pt,2) - 3s3p3d(ZD3,,2) (221.124) d 'f
3s3p2(Zp3 2) - 3s3p3d(2Ds,2) 223.220
3s3p2(eD3.,2) - 3s3p3d(2Fs,2) 224.354
3s3p(3po) - 3s3d(3D ~ ) 224.754 b
3s3pe(2Ds,.2) - 3s3p3d(2Fv,a) (225.481) < f
3s3p2(2St,,2) - 3s3p3d(2P3.,2) 226.040
3s3p2(4p3,,2) - 3s3p3d(4P3:2) 227.445
3p2(tD2) - 3p3d(3D3) 229.744 b
3p2(3pI ) - 3p3d(3D2) 231.47 c
3pZ(3P0) - 3p3d(3pl) 233.46 c
3s23d(2D3..2) - 3s3p3d(2D3 z) 235.776
3s3p2(2Pk,2) - 3s3p3d(2pv2) 237.369
3s3p2( 4P ..2)- 3p3(4S3.~) 238.699f
3p2(3p1) _ 3p3d(3p2) 238.708 b, r, g
3s23d(2D3:2) - 3s3p3d(zPt .~) 239.055
3p2(3P2) - 3p3d(3p3) 242.100 b. f
3s3pZ(4p5:2)- 3s3p3d(4Ds:2) 249.374
3s3p(3p i ) - 3p2(3p~) 251.949 b
3s3p2(2Ds,,2)- 3s3p3d(2D~.2) 257.234
3sZ3p(2p1,2)- 3s3p2(2p3.2) 294.758 ~-
3s23p(2Ds:2)- 3p3(2Ds,2) 313.724
3s3d(3D3) - 3p3d(3D3) 319.047 b, f
3s3d(ID2)- 3p3d(~F3) 332.854 b
3s3p=(2Ds:2) - 3p3(2Ds,~) 358.681
3s23p(2p3:2)- 3s=3d(2Ds,2) 359.369 r
3s23p(2P3:2)- 3s23d(2D3,2) 359.660
3s23p( Zp3,9_ ) - 3s3p2(2pt,,2 ) 365.628 r
3s3p2(=p3,,2) - 3s3p3d(2pl 2) 372.619 e
3s3p2(2Sl 2 ) - 3s3p3d(zP3,2) 374.100 r
3s3p=(4p1..2 ) - 3p3(4S3,2) 379.780 f
3s3p2(2Ds:2) _ 3p3(=p3. a ) 399.797 r, h
3s3p(3p2) - 3p2(3p~ ) 515-56 c
202.05
215.93
216.88
218.19
220.86
221.12
223.19
224.34
224.74
225.49
226.03
227.47
229.73
231.45
233.45
235.81
237.34
238.71
239.03
242.07
249.38
251.95
257.26
294.72
313.74
319.03
332.79
358.67
359.37
359.65
365.62
372.64
374.10
379.80
399.83
515.60
Fe Xllt
S x I
SiVlll; Fe x t l I
SxL~
Ar x v
S i x
S x n
HetI
H e l I
He lI
Fe x x I
FeXi l l i
M g v u t
N i x v ; M g v l l
A I •
F e x i ~
N e v
Fe Xltl
Ne V I
He~
a Obse rved wavelengths from RL, except as footnoted. b Chur i lov et al. (1985).
c Li tzen and Redfors (1987). a F r o m RL levels; MB line 16, t rans i t ion 9-35 .
e F rom RL levels; t rans i t ion 7-29. r Not present in B C F D quie t -Sun spectrum.
g Identified by Churi lov et al. (1985). h Two wavelength-coinc ident t ransi t ions .
Widing, Sandl in and Cowan (1971).
Fexl,, AND OTHER LINE IDENTIFICAFIONS IN THE SOLAR EUV SPECTRUM 349
Ne
TABLE vI
sequence 2p53s(3_Po) level values and predicted forbidden 2p53s ( 3 P o --~ 3P I ) wavelengths% from Buchet et aL (1987)
Ion E(3Po) 2(3P0 . 3pi) (em ~) (A)
Artx 2044650 (8688 + 15) c Kx 2422735 (6332 + 3) ~ Cax[ 2832095 (4718 + 23) c Scxll (3272751) b (3588 ) b
TixHI 3745200 (2774 + 12) c Vxw (4248722) b (2192.) b Crxv 4784450 1749.6 e Mnxv! (5351766) b (1417 .5 ) b
Fexvn 5951500 d 1153.20 f Nixlx (7247700) b (799.4) b
a This table replaces Table 4 of Kastner (1983). b Theoretical values. c Unobserved. a Deduced here from the observed Fexvr~ wavelength. e Tentative assignment; see text. f Feldman, Doschek, and Seely (1985).
abundances of the stable neon-like ions, none of the other lines in ArIx , Cax t , Tix~II,
or Ni x Ix have been observed in solar spectra. At the same time observed coronal lines
at 4566.2, 1696.26, 1452.68, and 1189.52 3~ (Jefferies, Orrall, and Zirker, 1971 ; Sandlin,
Brueckner, and Tousey, 1977; Sandlin and Tousey, 1979), which SK initially linked to
the neon-like lines, have not yet been identified.
4. Discussion of Individual Lines
Some entries in the tables deserve special comment:
(1) Blaha (1971) included two F e x I v transitions, MB line 9 [transition 6 -31] and
MB line 17 [transition 7-31] as being sufficiently intense to be observable in the Sun,
particularly the former. The new laboratory results of RL show that these lines have the
wavelengths 191.815 and 192.638,~, respectively, replacing longer wavelengths
suggested by SK. As one of the observational anomalies found in this survey, the latter
line is present in the B C F D spectrum (Table I II) but the former line is absent.
(2) Blaha's line 16 [transition 9-35] is observed by RE to lie at wavelength
221.124 A. It is not present in the B C F D spectrum, but is found in the flare spectrum
(Table V). The latter assignment replaces A r x v which is less probable on the basis of abundance.
(3) Both of the weaker Blaha transitions 7-38 [MB line 21] and 7-32 [MB line 25]
are present in the B C F D spectrum, at wavelengths 197.776 and 193.746 ]~, respectively
(Table lII). This makes the absence of t rans i t ion6-31 [MB line9], expected at
191,815 f \ in the same wavelength range, still more anomalous. The former transition
350 s.o. KASTNER
is associated here with the BCFD line 197.847 ~, despite the difference of 71 mA, since the observed line is wide.
(4) Blaha's line 18 [transition 10-36] at the RL wavelength 223.220 A_ is associated
in Table III with the BCFD line 223.202 k,. Since the observed line is wide, the transition 5-28, at a wavelength 223.261 A, not directly observed by RL but obtained
from the RL levels, is also likely to be a component.
(5) The entry footnoted 'e' in Table IV compares an F e x v line with a solar line
assigned by BCFD to Alx, despite a wavelength difference of about 70 mA. This is included because of the low abundance of aluminum and the expected appreciable
intensity of the Fe xv singlet transition t F 3 --* ~D 2. The rather large wavelength dis- crepancy makes the assignment tentative.
(6) Two transitions are possible from the F e x w level 3s3p2(2D3/2)=_ level 6: the
transitions 6-1 [MB line 4] at 334.171 A and 6 -2 expected at 356.640 ]~. The former is present in the solar spectrum, but the latter is absent. The calculation of Kastner
(1985) indicates that this line should be of greater intensity than MB line 14 at
289.160 A, so that it would have been expected to be observable in the flare spectrum.
(7) Several of the Fe, Ni, and Cr assignments in the tables are paired with previous
identifications ofNe, Mg, Si, or S ion lines. This raises a question of whether the heavier
and lighter groups of ions both contribute to the spectra, or whether only one group is mainly responsible. An answer to this question may require high-resolution wavelength
determinations and/or line profiles which could distinguish between the two possibili- ties; profiles should be narrower for the heavier ions.
The question of validating identifications has often been a difficult one which can take
much time to resolve, or remain unresolved. In later work calculations of expected line intensities will be carried out, and use will be made of a modification of a branching ratio
calibration technique introduced by Neupert and Kastner (1983), to test some of the presently suggested assignments.
Acknowledgements
I wish to thank W. E. Behring for a critical reading of the manuscript. A referee's comments have contributed to the clarity of the paper.
References
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Fexw 4.ND OTHER LINE IDENTIEICATIONS IN THE SOLAR EUV SPECTRUM 351
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