T H E O R E T I C A L E M I S S I O N L I N E S T R E N G T H S F O R NeVII

C O M P A R E D T O E U V S O L A R O B S E R V A T I O N S

F. P. KEENAN

Department of Pure and Applied Physics, The Queen's University of Belfast, Belfast BT7 INN, N. Ireland

(Received 29 August, 1990)

Abstract. Theoretical electron-temperature-sensitive NevH emission line ratios, calculated using accurate R-matrix electron impact excitation rates, are presented for R 1 = I(895.2]~)/I(465.2]~), R 2 = I(561.7 A.)/I(465.2 •) and R 3 = I(564.5 N)/I(465.2 ~). A comparison of these with observational data for several solar features obtained with the Harvard S-055 spectrometer on board Skylab reveals good agreement between theory and experiment. This provides observational support for the accuracy of the atomic physics adopted in the calculations, and the methods employed in the derivation of the theoretical diagnostics.

1. Introduction

Emission lines arising from An = 0 (2-2) transitions in Be-like ions are frequently detected in the spectra of the solar transition region and corona (Widing and Feldman,

1989; Sandlin et al., 1986). These transitions may be used to infer the electron density

(Ne) and temperature (Te) of the emitting plasma through diagnostic line ratios, although to determine these reliably requires the adoption of accurate atomic data in the calcu- lations, especially for electron impact excitation rates and f-values (Dufton and Kingston, 1981).

For several years we have been involved in an extensive series of Be-like ion N e- and

Te-diagnostic calculations for solar plasmas (see Keenan, 1990, and references therein),

using electron excitation rates either determined with the R-matrix code (Berrington

et al., 1977, 1985; Berrington, 1985; Dufton, Kingston, and Scott, 1983), or interpolated from these (Keenan et al., 1986; Keenan, 1988). Very recently, Keenan, McCann, and

Widing (1990) have used the atomic data of Berrington etal. (1985) for Nevl I to evaluate relative populations and, hence, emission line strengths for this ion, which incorporated several improvements over the previous calculations of Dufton, Doyle,

and Kingston (1979), including more accurate electron impact excitation rates and inclusion of reliable proton rates for transitions among the 2s2p 3_p levels (see Keenan et al., 1990, for more details). In this paper we use the Keenan et al. (1990) results to

derive theoretical line ratios, and compare these with EUV solar observations from Skylab to investigate their usefulness as Te-diagnostics.

2. Theoretical Ratios

The model ion adopted for N e v n has been discussed in detail by Keenan et al. (1990). Briefly, the six energetically lowest LS states were included in the calculations, namely 2s a 1S, 2s2p 3p, 1p; 2p2 3p, 1D, and 1S, making a total of ten levels when the fine

Solar Physics 131: 291-296, 1991. �9 1991 Kluwer Academic Publishers. Printed in Belgium.

292 V. P. KEENAN

structure splitting in the triplet terms is included. Only collisional excitation and

de-excitation by electrons and protons (the latter in the case of transitions among 2s2p 3p), and spontaneous radiative de-excitation processes were considered, and the

plasma was assumed to be optically thin. Further details may be found in Keenan et al.

(1990). In Figures 1-3 we plot the theoretical ratios

R 1 = I(2s 2 i S - 2s2p 3p1)/I(2s2 i s - 2s2p 1 p ) ,

R2 = I ( [2 -2 ] + [1-1]) / I (2s 2 1S - 2s2p 1p),

and

R3 = I ( [ 2 - 1 ] ) / I ( 2 s 2 1S - 2 s 2 p 1tO)

as a function of electron temperature at a density of N e = l01 ~ cm - 3, where we have

used [J' - J ' ] to denote a transition of the type 2s2p 3p, _ 2pZ 3pj,. These ratios are

insensitive to variations in the electron density when Ne-< 3 X 10 ~1 cm -3, and an

o -#

r ' r "

7

{5-

5 -

/~-

3 -

2 -

1 -

I I I I

I I I I 5.L 5.7 6.0 6.3

t0g Te Fig. 1. The theoretical NevlI emission line ratio

R 1 = I ( 2 S 2 1S - 2s2p 3P1)/1(2S2 'S - 2s2p 1p) = 1(895 .2 ,&)/I(465.2 • ) ,

where I is in energy units, plotted as a function of electron temperature at an electron density of N~= 10 l l c m 3.

E M I S S I O N L I N E S T R E N G T H S F O R N e V I I C O M P A R E D T O E U V S O L A R O B S E R V A T I O N S 293

. - - . -__ . .

_

3 -

2 -

l I I I

I I t I 5.4 5.7 15.0 6.3

tog Te Fig. 2. The theoretical NevtI emission line ratio

R 2 = l(2s2p 3 P 2 - 2 p 2 3 P 2 + 2s2p 3 P 1 - 2p 2 3p1)/1(2s2 1S - 2s2p l p ) = I(561.7 ~)/I(465.2 ~ ) ,

where I is in energy units, plotted as a function of electron temperature at an electron density of N e = 1011 cm -3.

inspection of the figures shows that R 1 - R 3 are, hence, potentially very useful

T~-diagnostics for all but the highest density solar features. For example, R~ varies by

a factor of 4.6 between log T e = 5.4 and 6.3, while R 2 changes by a factor of 2.3 over the same temperature interval.

3. Resu l t s and D i s c u s s i o n

The2s 2 IS - - 2 s 2 p 3p1, [2-1], [2-2] + [1-1] and2s a 1S - 2 s 2 p ~Ptransitions in NevII

have been observed at wavelengths of 895.2 A, 564.5 A, 561.7 A,, and 465.2 A, respec-

tively, in solar emission line spectra obtained with the Harvard S-055 spectrometer on board S k y l a b (Doyle, 1983). This instrument, which covered the wavelength region 280-1350 A, observed a spatial area of 5 x 5 arc sec with a spectral resolution of approximately 1.6 A (FWHM) using an integration time of 0.04 s and a step length of 0.2112 A. It is discussed in detail by Reeves, Huber, and Timothy (1977) and Reeves et al. (1977).

In Table I we summarise the observed values of R 1 = I(895.2 A,)/I(465.2 A_),

294 F.P. KEENAN

1.3

1.2

1.1 O

. m

# 1.0 %

0.9

0.8

0.7

0.6

I I I I

I 1 I I 5J., 5.7 6.0 6.3

log Te Fig. 3. The theoretical NevI] line ratio

R 3 = I ( 2 s 2 p 3P 2 - 2p 2 3P1)/1(2s2 18 - 2 s2p ip) = I(564.5 A)/I(465.2 A) ,

where I is in energy units, plotted as a function of electron temperature at an electron density of Ne = 1011 cm -3.

R 2 = I(561.7 A)/I(465.2 A), and R 3 = 1(564.5 *)/1(465.2 A) for several solar features, along with the sources of the data. All these events have electron densities

N e < 3 x 1011 cm 3 (see individual references in Table I), so that the R 1 - R 3 line ratios are only sensitive to variations in electron temperature. Also shown in the table are the electron temperatures derived from these ratios using the calculations in Figures 1-3. An inspection of the table reveals that the values of T e determined from R 1 - R 3 are generally compatible, with discrepancies of typically 0.3 dex. This corresponds to only a ,-~20% change in R 2 or R3, which is well within the estimated errors in the observational data ( ~ 3 0 % ; Noyes e t a l . , 1985). Furthermore, the derived electron temperatures are in excellent agreement with that of maximum N e v n fractional abun- dance in ionisation equilibrium, log Tma x = 5.7 (Arnaud and Rothenflug, 1985), with differences averaging only 0.2 dex. This provides observational support for the accuracy of the atomic data adopted in the calculations, and the methods employed in the determination of the theoretical line ratios.

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296 F.P. KEENAN

4. Conclusions

There are two principal solutions:

(1) Theoretical Nevl I emission line ratios R 1 = I ( 8 9 5 . 2 A ) / I ( 4 6 5 . 2 A ) , R2 =

= I(561.7 A)/I(465.2 A), and R 3 = 1(564.5 A)/I(465.2 A), derived using accurate

electron impact excitation rates calculated with the R-matrix code by Berrington et al.

(1985), are found to vary significantly with electron temperature, but to be insensitive

to changes in the electron density when N~ < 3 x 1011 c m - 3. Hence, they are potentially

useful Te-diagnostics for all but the highest density solar features.

(2) Electron temperatures derived from the observed value of R 1 - R 3 for several

solar features obtained with the Harvard S-055 spectrometer on board Skylab are

compatible, with discrepancies of typically 0.3 dex. Furthermore, they are in good

agreement with the temperature of maximum N e v n fractional abundance in ionisation

equilibrium, log Tma • = 5.7 (Arnaud and Rothenflug, 1985), with differences averaging

only 0.2 dex. This provides experimental support for the accuracy of the atomic data

adopted in the calculations, and the methods employed in the derivation of the theoreti-

cal diagnostics.

Acknowledgement

I would like to thank Prof. A. E. Kingston for his continued interest in this work.

References

Arnaud, M. and Rothenflug, R.: 1985, Astron. Astrophys. Suppl. 60, 425. Berrington, K. A.: 1985, J. Phys. B18, L395. Berrington, K. A., Burke, P. G., Dufton, P. L., and Kingston, A. E.: 1977, J. Phys. B10, 1465. Berrington, K. A., Burke, P. G., Dufton, P. L., and Kingston, A. E.: 1985, Atomic Data Nucl. Data Tables

33, 195. Doyle, J. G.: 1983, Solar Phys. 89, 115. Doyle, J. G., Mason, H. E., and Vernazza, J. E.: Astron. Astrophys. 150, 69. Dufton, P. L. and Kingston, A. E.: 1981, Adv. Atom. Molec. Phys. 17, 355. Dufton, P. L., Doyle, J. G., and Kingston, A. E.: 1979, Astron. Astrophys. 78, 318. Dufton, P. L., Kingston, A. E., and Scott, N. S.: 1983, J. Phys. B16, 3053. Keenan, F. P.: 1988, Phys. Scripta 37, 57. Keenan, F. P.: 1990, Solar Phys. 126, 311. Keenan, F. P., Berrington, K. A., Burke, P. G., Dufton, P. L., and Kingston, A. E.: 1986, Phys. Scripta 34,

216. Keenan, F. P., McCann, S. M., and Widing, K. G.: 1990, Astrophys. J. (in press). Noyes, R. W., Raymond, J. C., Doyle, J. G., and Kingston, A. E.: 1985, Astrophys. J. 297, 805. Reeves, E. M., Huber, M. C. E., and Timothy, J. G.: 1977, AppL Opt. 16, 837. Reeves, E. M., Timothy, J. G., Huber, M. C. E., and Withbroe, G. L.: 1977, Appl. Opt. 16, 849. Sandlin, G. D., Bartoe, J.-D. F., Tousey, R., and Van Hoosier, M. E.: 1986, Astrophys. J. Suppl. 61,801. Vernazza, J. E. and Reeves, E. M.: 1978, Astrophys. J. Suppl. 37, 485. Widing, K. G. and Feldman, U.: 1989, Astrophys. J. 334, 1046.