P R O P O S E D M g v , Fe x I v A N D F O R B I D D E N Fe xvI I LINE

I D E N T I F I C A T I O N S IN THE EUV SOLAR S P E C T R U M

(Research Note)

S O. K A S T N E R

Laboratory for Astronomy and Solar Physics, NASA-Goddard Space Flight Center, Greenbelt, MD 20771, U.S.A.

(Received 21 December, 1981; in revised form 22 March, 1982)

Abstract. Some line identifications in solar disk and limb spectra are proposed on the basis of recent laboratory and theoretical results reported in the literature, including allowed lines of Mg I and Fe XlV in the EUV spectrum and an expected forbidden line of FexvI1 near 1190 ~.

1. Introduction

In the high-resolution spectra obtained by Behring et al. (1972, 1976) and in the spectra obtained by the Skylab ATM experiment, a number of lines remain unidentified. On the other hand, certain transitions can be expected on theoretical grounds to be of sufficient intensity to be present in these spectra, in particular some allowed lines of Mg v and Fe xlv which have not yet been identified. The following assignments are suggested on the basis of wavelength and intensity considerations.

2. Allowed Lines of Mg v

Kaufman and Artru (1980) have given an improved list of Mgv lines in the extreme ultraviolet. Several of these have the ground levels 2p 4 (3P2, 1) as their lower level and can t~refore be excited under solar conditions. Of these, Behring et al. (1972) have observed the lines 114.052 ]~ and 132.492 ,~. The unidentified line 146.083 A in their spectrum can now be classified as 2pn(3p2) - 2p3(4S)3s(3S1) on the basis of exact wavelength agreement with a line in the list of Kaufman and Artru. Other lines at 121.644 A, 113.699 A and 126.280 A appear to be present in the photoelectric spectrum of Manson (1972), and the line 137.412 ]~ is to be expected under favorable observing conditions. The lines at 276.582 A and 353.092 A should also be observable in a more sensitive spectral scan.

3. Allowed Lines of Fe xIv

Blaha (1971) made an extensive study of the Fexw spectrum, obtaining expected intensities of all important EUV lines as well as a number of weaker lines. Table I compares his predicted intensities with photometric intensities observed by Malinovsky

Solar Physics 81 (1982) 59-61. 0038-0938/82/0811-0059500.45. Copyright �9 1982 by D. Reidel Publishing Co., Dordrecht, Holland, and Boston, U.S.A.

60 S. O. KASTNER

TABLE I

Predicted and observed intensities of weaker Fe xlv EUV lines [- log (12/1211.3) ]

Observed

Predicted" MH b HCH c BCFD a

252.20 0.94 0.89 - 0.57 242.215 e 1.16 1.3 f - 0.75 219.12 0.78 0.73 0.65 0.48 353.83 1.29 - - 0.70 289.16 1.64 1.45 - 1.05 243.42 e 1.84 - - > 1.65

a Blaha (1971); electron density N e = 1 x 109 cm -3. b Malinovsky and Heroux (1973). c Heroux etal. (1974). d Behring et al. (1976). e Present identifications. f Present estimate of intensity in MH spectrum.

and Heroux (1973) (MH) and Heroux et aL (1974) ( H C H ) for the weaker F e x l v lines.

The intensities of all the lines so far observed are in agreement with Blaha's predictions

so that one may rely on his values for the lines not yet observed or identified. The visual

intensity estimates o f Behring et al. (1976) (BCFD) have been included also in the last

column of Table I to show that they are in agreement with the photometric observations

and the predictions with respect to the ordering of intensities.

Of the lines as yet unidentified, one pair is expected to originate from the unusual level

3s3p3d(2F'5/2). The pair consists of a relatively strong transition to the 3s3p2(2D3/2) level and a weaker transition to the 3s 3p2(2Ds/2) level. The former line is expected to

be actually stronger than the observed line 289.16 A, and a search for the pair was

therefore undertaken in the high-resolution spectrum of B CFD. Two lines at 242.215

and 243.42 A appear to satisfy the necessary conditions. As may be seen in Table I, their

relative intensities are in good agreement with the expected intensities. The stronger line

in fact was noted to be present also in the spectrum of MH, enabling a quantitative

intensity value to be obtained. Their separation is 2045 + 50 c m - 1, in reasonable agree-

ment with the known 3s3p2(203/2 - 205/2) separation o f 2224 c m - 1 considering the

uncertainty of the quoted error estimate for the weaker line. The position of the 3s3p3d(2F'5/2) level is located at 712 105 + 100 c m - ~ on the basis o f this identification.

4. A Forbidden Line of Fe xvn

Of the four levels o f the 2p 53s configuration in neonlike Fe xvII, three (~P1, 3pI, 3P2) can decay to the ground level 2p6(1So) with high or relatively high probability but the J = 0 level decay is strictly forbidden. The 3P o level is therefore metastable and will be

highly populated in low density plasmas. Some time ago Gars tang (1969) calculated

LINE IDENTIFICATIONS IN THE EUV SOLAR SPECTRUM 61

magnetic dipole and electric quadrupole transition probabilities within this configura- tion, obtaining a relatively high magnetic dipole transition probability of 1.6 x 104 s- 1 for the transition 3P o ---, 3P 1 . From the results of a recent theoretical study (Bogdanovich etal., 1978) the position of the 3P o level can be estimated as 5950000 + 2000 cm -1, giving the expected wavelength of this transition as 1170 + 35 ,~. Since FexvII is present to some extent in the quiet corona as evidenced by quiet Sun lines of Fe xvIII at 93.94 and 103.96 A, the transition can be expected to be present, though weak, in the coronal spectrum. Sandlin and Tousey (1979) report, in a list of weak coronal lines observed above the solar limb by the ATM experiment, a line at 1189.52 ]~. It is suggested that this may be the 2p53s(3po ~ 3P1) transition of Fexvu. Its intensity is of the order of that of another iron forbidden line in the same list. If this identification were verified the location of the 3P o level would be established at 5 948 828 + 500 cm- 1. A possibility for verification would be to identify the corresponding line in NixIx, in low density laboratory plasmas. The latter wavelength is to be expected at 820 + 30 ~, if the above coronal line is indeed due to FexvII.

References

Behring, W. E., Cohen, L., and Feldman, U.: 1972, Astrophys. J. 175, 493. Behring, W. E., Cohen, L., Feldman, U., and Doschek, G. A.: 1976, Astrophys. J. 203, 521. Blaha, M.: 1971, Solar Phys. 17, 99. Bogdanovich, P. O., Rudzikas, Z. B., Safronova, V. I., and Shadzhyuvene, S. D.: 1978, Opt. Spectrosc.

(USSR) 44, 618. Garstang, R. H.: 1969, Publ. Astron. Soc. Pacific 8I, 488. Heroux, L., Cohen, M., and Higgins, J. E.: 1974, J. Geophys. Res. 79, 5237. Kaufmann, V. and Artru, M.-C.: 1980, J. Opt. Soc. Am. 70, 1135. Malinovsky, M. and Heroux, L.: 1973, Astrophys. J. 181, 1009. Manson, J. E.: 1972, Solar Phys. 27, 107. Sandlin, G. D. and Tousey, R.: 1979, Astrophys. J. 227, L107.