THE ASTROPHYSICAL JOURNAL, 544 : 508È521, 2000 November 20
( 2000. The American Astronomical Society. All rights reserved. Printed in U.S.A.
IDENTIFICATION OF SPECTRAL LINES IN THE 500È1600 AŽ WAVELENGTH RANGE OF HIGHLY
IONIZED Ne, Na, Mg, Ar, K, Ca, Ti, Cr, Mn, Fe, Co, AND Ni EMITTED BY FLARES (T º 3 ] 106 K)
e
AND THEIR POTENTIAL USE IN PLASMA DIAGNOSTICS
U. FELDMAN,1 W. CURDT,2 E. LANDI,2 AND K. WILHELM2
Received 2000 April 6 ; accepted 2000 June 29
ABSTRACT
On 1999 May 9 the Solar Ultraviolet Measurements of Emitted Radiation (SUMER) spectrometer on
the Solar and Heliospheric Observatory (SOHO) recorded spectra from a high-temperature region located
in the solar corona above the west limb. These spectra contain lines from rather less-abundant elements
in solar plasmas. In this paper we present identiÐcations of the high-temperature (T º 3 ] 106 K) Ne,
e
Na, Mg, Ar, K, Ca, Ti, Cr, Mn, Fe, Co, and Ni lines that were detected in the 500È1600 AŽ spectral range
of SUMER. In addition, accurate wavelength measurements have been obtained with uncertainties
varying between 0.015 and 0.040 AŽ (1 p). Making use of the newly measured wavelengths, we derive
energy levels in the ground conÐguration of a number of highly charged ions. We present intensity ratio
calculations of lines in the SUMER range that could be used to measure electron densities in hightemperature solar plasmas. We also provide emissivities for Ca XIIIÈCa XV and Fe XVIIIÈFe XXIII lines
that could be used to determine emission measures and electron temperatures of high-temperature
plasmas. We discuss a method for measuring elemental abundance variations in high-temperature solar
plasmas using lines presented in the paper. A list of spectral lines spanning the 300È30000 AŽ wavelength
range and their branching ratios that are suitable for efficiency calibration of space-borne spectrographs
is provided.
Subject headings : line : identiÐcation È Sun : abundances È Sun : corona È Sun : Ñares È
Sun : UV radiation
1.
INTRODUCTION
to determine the electron densities, electron temperatures,
emission measure distributions, mass motions, and the state
of ionization equilibrium of the plasmas. The lines emitted
by ions from many di†erent elements can also be used to
determine the relative elemental abundances in hot astrophysical plasmas.
The scope of this paper is to derive analysis methods for
plasma diagnostics. A future paper will apply these methods
to the 1999 May 9 Ñare. In ° 2 of the paper we brieÑy discuss
the SUMER instrument, and in ° 3 data acquisition and
reduction are described. The identiÐcations of the hightemperature lines, their wavelengths, and the energy levels
inferred from them are discussed in ° 4. In ° 5 we discuss line
intensity ratios, which could be used to derive electron densities. Section 6 deals with emission measure determinations, and ° 7 with electron temperature determinations
and the state of ionization equilibrium using hightemperature lines emitted in the 500È1600 AŽ range. Elemental abundance determinations are discussed in ° 8, and in ° 9
we explore the possibility of using sets of the forbidden lines
for efficiency calibration of spectrometers.
On 1999 May 9 the Solar Ultraviolet Measurements of
Emitted Radiation (SUMER) spectrograph on the Solar
and Heliospheric Observatory (SOHO) recorded farultraviolet spectra from a bright region of the corona
located above the west limb. During the 6 hours and 40
minutes of observations the instrument obtained two complete spectra spanning the entire SUMER wavelength range
from 500 to 1610 AŽ . The spectral recordings were done in
sections of 43 AŽ , where for each section the start wavelength
was advanced by B 20 AŽ relative to the preceding section.
The integration time at each 43 AŽ section was 300 s. The
tenth Geostationary Operational Environmental Satellite
(GOES 10) records show the eruption of several Ñares of
di†erent sizes on the Sun from about 12 : 00 UT. In a companion paper based on images from SUMER, EIT, Y ohkohSXT, and the MICA coronagraph data, D. E. Innes (2000,
private communication) have demonstrated that the Ñare
onset occurred on the sunward side of the SUMER slit
position from where it rapidly expanded to the site observed
by the spectrometer. While the recording of the Ðrst spectra
was underway, a M7.6 Ñare erupted in the observed active
region. As a result, the spectra that were acquired during
this Ñare contain the hottest high-resolution astrophysical
spectral recordings ever made in the 500È1600 AŽ range.
Although resonance lines emitted from the hightemperature (T º 3 ] 106 K) astrophysical plasmas have
e
wavelengths shorter
than 500 AŽ , the SUMER instrument
could record many bright forbidden lines also emitted by
Ñare plasmas. The intensity of many of the lines can be used
2.
THE SUMER INSTRUMENT
The SUMER instrument is composed of a telescope and
a spectrometer capable of producing stigmatic highresolution spectra. The telescope consists of an o†-axis
parabolic mirror that can image any region within a
64 ] 64 arcmin2 Ðeld of view centered on the Sun. (The
solar diameter as seen from SOHO is B32@.) The spectrometer entrance slit is placed in the focal plane of the telescope. The spectrometer consists of the entrance slit, an
o†-axis parabolic mirror which collimates the light leaving
the slit, a Ñat mirror which deÑects the light onto a concave
grating in a Wadsworth conÐguration, and two imaging
detectors. The two detectors, aligned with the grating focal
plane, can be used alternatively to collect stigmatic images
1 E. O. Hulbert Center for Space Research, Naval Research Laboratory,
Washington, DC 20375-5352.
2 Max-Planck-Institut fuŽr Aeronomie (MPAE), Max-Planck-Strae 2,
D-37191 Katlenburg-Lindau, Germany.
508
SPECTRAL LINES AND THEIR POTENTIAL USE IN PLASMA DIAGNOSTICS
of the slit. Detector A, which was used during our observations, covers the 780È1610 AŽ range in Ðrst order. Secondorder lines are superposed on the Ðrst-order spectrum. The
detector has an array of 1024 (spectral) ] 360 (spatial)
pixels each on average 26.5 ] 26.5 km2 in size and covers a
spectral range of 43 AŽ in Ðrst order. The angular scale of a
pixel at 800 AŽ is 1A. 03, while at 1600 AŽ it is 0A. 95. In the
Ðrst-order spectrum a pixel corresponds to 45.0 mAŽ at 800
AŽ and 41.8 mAŽ at 1600 AŽ . With SOHO in its nominal attitude the SUMER instrument is oriented with its slit aligned
along the north-south direction. During the observations a
1 ] 300 arcsec2 slit was used.
The SUMER optics are made of silicon carbide, which is
a fairly good reÑector for radiation with wavelengths longer
than 500 AŽ . Each photon recorded by SUMER is reÑected
by three normal incidence and one grazing incidence silicon
carbide surfaces before reaching the detector. The secondorder spectrum superimposed on the Ðrst order should
cover the 390È805 AŽ range. However, as the radiation wavelength approaches 500 AŽ from the long wavelength side, the
reÑective properties of the SUMER optics diminish rapidly.
The detector photocathode surface is divided into a
number of sections. Two small sections of 55 pixels at both
extremes of the detector are covered with a grid that transmits only 10% of the incident light. The next 210 pixels on
both sides of the microchannel plate detector are uncoated,
while the center 490 pixels are coated with potassium
bromide (KBr). The efficiency of the detector section coated
with KBr is higher than the efficiency of the uncoated parts
over most of the wavelength range. Therefore, a comparison
of the intensity of a particular line recorded on the KBr with
its intensity obtained from the uncoated part of the detector
can reveal unambiguously if a line is seen in Ðrst or second
order. The details of the SUMER instrument and its modes
of operation are described in Wilhelm et al. (1995). Wilhelm
et al. (1997) provide an account of the actual performance of
SUMER under operational conditions.
3.
DATA ACQUISITION AND REDUCTION
On May 9, SUMER observed an o†-limb postÑare site
above the active region NOAA 8537, which was approaching the solar west limb. SUMER was in the so-called reference spectrum mode, in which the instrument scanned with
some overlap the entire wavelength range in D 43 AŽ sections. This operation lasted from 15 : 24 to 18 : 38 UT and
was repeated from 18 : 44 to 21 : 57 UT. A Ñare of size C5.1
occurred at 16 : 09 UT, when the instrument recorded the
spectrum around 980 AŽ . At about 17 : 53 UT, a solar Ñare of
size M7.6 erupted. The X-ray Ñux from the Ñare, as monitored by the GOES 8 satellite, peaked at about 18 : 07 UT,
and since then it began its decay. By 19 : 00 UT the X-ray
Ñux had dropped by over 1 order of magnitude from its
maximum, and by 21 : 00 UT it dropped another factor of 2
to 3 until it could not be distinguished from the solar X-ray
background. At Ñare onset the detector recorded spectra
near 1395 AŽ while advancing toward longer wavelengths. At
the time of peak X-ray Ñux the detector recorded the spectrum around 1450 AŽ .
The bright region in the corona, which the SUMER slit
was pointed at during the 1999 May 9 observations is
shown in Figure 1. The observed region was the site of
continuous Ñare activity. Therefore even prior to the onset
of the main Ñare emission from hot plasma was detected,
509
although some of the hottest lines are only observed later
during the main phase of the M Ñare.
Standard procedures have been applied to the basic data
processing, i.e., decompression, Ñat-Ðeld correction, and
geometrical distortion corrections. The pixel-to-wavelength
relation in SUMER spectra is wavelength dependent and
very sensitive to mechanism slack. Therefore it needs to be
calibrated for each exposure. This is achieved by a correlation of the entire 43 AŽ window with all known reference
lines in this window. Assuming that the dispersion is known
very accurately from the optical design, this correlation
leads to a constant o†set for each exposure. The method
leads to a pixel-to-wavelength calibration with systematic
uncertainties from 10 to 50 mAŽ , depending on the quality of
the wavelength standards in the exposure. Another error
source comes from the determination of the line centroids.
Centroiding can be difficult for blends or non-Gaussian line
proÐles, but normally the uncertainty is better than 4 mAŽ
for isolated lines. The sum of both contributions yields the
uncertainty of each individual measurement.
Most of the lines have been observed twice during each of
the spectral scans, so that in most cases four measurements
are available. Their weighted averages are listed in Table 1.
4.
HIGH-TEMPERATURE SPECTRAL LINES IN THE SUMER
RANGE AND INFERRED ENERGY LEVELS
The spectra recorded on 1999 May 9 contain a large
number of spectral lines emitted by hot (T º 3 ] 106 K)
e
plasmas. Some of the recorded lines were
previously
observed in solar spectra (e.g., Doschek et al. 1975 ; Sandlin,
Brueckner & Tousey 1977 ; Feldman & Doschek 1991 ;
Feldman et al. 1998a ; Kucera et al. 2000). A number of
additional lines were previously observed in spectra emitted
by tokamak plasmas and by beam-foil sources (e.g., Denne
& Hinnov 1984 ; Hinnov & Suckewer 1980 ; Hinnov et al.
1982 ; Peacock, Stamp, & Silver 1984). A list of the hightemperature lines detected in the 1999 May 9 spectra is
given in Table 1. Some of the Ñare lines have never been
identiÐed before. The table includes the line identiÐcations,
if available, their previously most accurately derived wave-
FIG. 1.ÈEIT images taken before and after the Ñare onset. The
SUMER slit position is superimposed on the Ðeld of view.
510
2s22p5 2P È2s22p5 2P
1@2
3@2
2s22p3 2P È2s22p3 2D
1@2
3@2
2p 2P È2p 3P
3@2
1@2
2s22p3 2D È2s22p3 4S
3@2
3@2
1s2p 3P È1s2s 3S
0
1
2s22p3 2D È2s22p3 4S
5@2
3@2
2s22p2 3P È2s22p2 3P
1
0
2s22p4 1D È2s22p4 3P
2
2
2s22p3 2D È2s22p3 4S
3@2
3@2
2s22p3 2D È2s22p3 4S
5@2
3@2
2s22p23p4D È2s22p23s 4P
7@2
5@2
2s22p3 2P È2s22p3 2D
3@2
3@2
2s22p5 2P È2s22p5 2P
1@2
3@2
2s22p3 2D È2s22p3 4S
3@2
3@2
1s2p 3P È1s2s 3S
2
1
2p 2p È2p 2P
1@2
3@2
2s22p3 2D È2s22p3 4S
5@2
3@2
3s23p2 1S È3s23p2 3P
0
1
3s23p3 2P È3s23p3 4S
3@2
3@2
1s2p 3P È1s2s 3S
0
1
2s22p3 2D È2s22p3 4S
3@2
3@2
Co XIX . . . . . . .
Fe XX . . . . . . . .
Fe XXII . . . . . .
Ti XVI . . . . . . .
Si XIII . . . . . . . .
Ca XIV . . . . . . .
Ni XXIII . . . . . .
Ti XV . . . . . . . .
Ca XIV . . . . . . .
K XIII . . . . . . . .
S X ...........
Ti XVI . . . . . . .
Fe XVIII . . . . . .
K XIII . . . . . . . .
Mg XI . . . . . . .
Mn XXI . . . . . .
Ar XII . . . . . . . .
Ni XV . . . . . . . .
Ni XIV . . . . . . .
Mg XI . . . . . . .
Ar XII . . . . . . . .
? .............
? .............
Fe XXIII . . . . . .
Ca XIV . . . . . . .
2s2p 3P È2s2p 3P
2
1
2s22p3 2P È2s22p3 4S
3@2
3@2
2s22p3 2D È2s22p3 4S
3@2
3@2
2s22p53s3P È2s22p53s 3P
0
1
1s2p 3P È1s2s 3S
2
1
Transition
Cr XVIII . . . . . .
Ni XIX . . . . . . .
Si XIII . . . . . . . .
Linea
TABLE 1
1054.585 ^ 0.015 (bl)
1062.450 ^ 0.040
1077.172 ^ 0.025
1079.415 ^ 0.025
1090.450 ^ 0.030/2 \ (545.225)
1005.862 ^ 0.020
1018.790 ^ 0.020
1033.041 ^ 0.015
1034.485 ^ 0.020
1043.282 ^ 0.020 (bl)
994.442 ^ 0.020
997.456 ^ 0.015
880.401 ^ 0.015
910.855 ^ 0.028
919.730 ^ 0.019
943.585 ^ 0.020
945.880 ^ 0.020
946.307 ^ 0.020
968.868 ^ 0.015 (bl)
974.850 ^ 0.020
820.019 ^ 0.025
821.706 ^ 0.030
845.570 ^ 0.020
861.799 ^ 0.020
878.690 ^ 0.040
793.160 ^ 0.021
794.605 ^ 0.025
814.715 ^ 0.025
Flare j
(AŽ )
1079.3 ^ 0.3
545.26^0.02
1018.87
1033.04
1034.48
1043.29 ^ 0.08
1043.314 ^ 0.08
1054.62
997.38 ^ 0.084
997.451 ^ 0.019
946.29
968.8 ^ 0.3
974.86 ^ 0.02
974.86^0.02
821.74^0.05
845.55 ^ 0.1
861.8 ^ 0.1
878.85
878.684 ^ 0.029
884.43
911.0
919.73 ^ 0.08
943.61^0.62
814.71 ^ 0.02
814.729 ^ 0.024
793.3 ^ 0.3
Previously
Measured j
(AŽ )
6.3
[6.2
[6.2
7.1
6.45
6.95
6.3
6.35
6.3
6.8c
6.35
6.8
6.45
7.05
6.5
6.45
6.35
[6.2
6.6
6.7
6.65
6.9
7.05
6.6
6.95c
6.6
B6.6
6.95c
log T b
e
(K)
HOT (T [ 2 ] 106 K) EMISSION LINES OBSERVED IN 1999 MAY 9 FLARE SPECTRA
e
Hinnov et al. 1982
Feldman et al. 1998a
Feldman et al. 1998a
Feldman et al. 1998a
Feldman et al. 1998a
Klein et al. 1985
Curdt et al. 2000
Feldman et al. 1998a
Klein et al. 1985
Curdt et al. 2000
Kink & EngrsoŽm 1999
E. Hinnov, unpublished 1986
Peacock et al. 1984
Feldman et al. 1998a
Kucera et al. 2000
Hinnov & Suckewer 1980
E. Hinnov, unpublished 1986
Howie et al. 1994
Curdt et al. 2000
Feldman et al. 1998a
Hinnov et al. 1982
Peacock et al. 1984
Feldman et al. 1998a
Howie et al. 1994
Curdt et al. 2000
Hinnov et al. 1982
References to Previous
Measurements
2
5
5
3
5
5
5
5
5
5
3
5
7
5
Cross Reference
to Other Tables
511
2s22p3 2P
Ca XIV . . . . . . . . .
? ................
Mn XVII . . . . . . . .
Fe XXI . . . . . . . . . .
Ni XIV . . . . . . . . . .
1s2p 3P È1s2s 3S
0
1
2s22p3 2P È2s22p3 2D
3@2
3@2
2s22p4 1S È2s22p4 3P
0
1
2s22p3 2P È2s22p3 4S
3@2
3@2
Ne IX . . . . . . . . . . .
Ca XIV . . . . . . . . .
Ca XIII . . . . . . . . .
Ar XII . . . . . . . . . .
? ................
? ................
Ar XIII . . . . . . . . . .
? ................
? ................
Cr XVIII] . . . . . .
2s22p3 2D È2s22p3 4S
5@2
3@2
2s22p2 1S È2s22p2 3P
0
1
2s22p4 1D È2s22p4 3P
2
2
3s23p4 1S È3s23p4 3P
0
1
2s22p 2P È2s22p 2P
3@2
1@2
1s2p 3P È1s2s 3S
2
1
2s22p4 1D È2s22p4 3P
2
2
2s22p5 2P È2s22p5 2P
1@2
3@2
2s22p2 1D È2s22p2 3P
2
1
3s23p3 2P È3s23p3 4S
1@2
3@2
K XII . . . . . . . . . . .
Ni XIII . . . . . . . . . .
Fe XIX . . . . . . . . . .
? ................
? ................
? ................
Cr XX . . . . . . . . . .
Ne IX . . . . . . . . . . .
2s22p53s 3P È2s22p53s 3P
0
1
Fe XVII . . . . . . . . .
È2s22p3 4S
3@2
2s22p2 1D È2s22p2 3P
2
1
2s22p2 1S È2s22p2 3P
0
1
1s2p 3P È1s2s 3S
2
1
2s22p4 3P È2s22p4 3P
1
2
2s22p4 1D È2s22p4 3P
2
2
2s22p3 2D È2s22p3 4S
5@2
3@2
Ca XV . . . . . . . . . .
Ca XV . . . . . . . . . .
Na X . . . . . . . . . . .
Fe XIX . . . . . . . . . .
Ca XIII . . . . . . . . .
Fe XX . . . . . . . . . .
1@2
Transition
Linea
1277.790 ^ 0.050 (bl)
1291.607 ^ 0.030
1297.398 ^ 0.030/2 \ (648.699)
1298.186 ^ 0.040/2 \ (649.093)
1308.146 ^ 0.030/2 \ (654.073)
1309.752 ^ 0.030/2 \ (654.876)
1313.344 ^ 0.020/2 \ (656.672)
1314.650 ^ 0.040/2 \ (657.325)
1320.402 ^ 0.025/2 \ (660.201)
1325.901 ^ 0.030/2 \ (662.951)
1256.484 ^ 0.025
1277.204 ^ 0.025
1184.467 ^ 0.030/2 \ (592.234)
1188.108 ^ 0.025
1192.978 ^ 0.040
1201.192 ^ 0.035
1205.728 ^ 0.025
1248.088 ^ 0.015
1159.706 ^ 0.030/2 \ (579.853)
1164.813 ^ 0.040
1167.760 ^ 0.015
1171.531/2 ^ 0.030 \ (585.766)
1174.656 ^ 0.020
1153.151 ^ 0.025
1098.484 ^ 0.040
1110.764 ^ 0.040/2 \ (555.382)
1111.754 ^ 0.030
1118.078 ^ 0.015
1133.756 ^ 0.020
1135.650 ^ 0.025/2 \ (567.825)
Flare j
(AŽ )
663.1 ^ 0.3
656.69^0.02
1277.23
1277.22^0.02
1277.71 ^ 0.02
1291.61^0.02
648.68^0.02
1205.9 ^ 0.3
1248.07 ^ 0.02
1248.097 ^ 0.014
585.8 ^ 0.3
1174.65
1174.65^0.02
592.234 ^ 0.006
1111.759 ^ 0.017
1118.060 ^ 0.01
1133.79
567.76
567.84^0.02
1153.20
1153.16
579.85
1098.44
Previously
Measured j
(AŽ )
TABLE 1ÈContinued
6.5c
6.45
6.35
6.3
[6.2
[6.2
6.4
[6.2
[6.2
6.75
6.3
6.25
6.8
[6.2
[6.2
[6.2
6.9
6.5c
6.45
[6.2
6.55
6.95
6.3
6.45
6.55
6.55
6.7c
6.8
6.35
6.9
log T b
e
(K)
Denne & Hinnov 1984
Feldman et al. 1998a
Sandlin et al. 1977
Feldman et al. 1998a
Brown et al. 1985
Feldman et al. 1998a
Feldman et al. 1998a
Hinnov et al. 1982
Beyer et al. 1986
Curdt et al. 2000
Hinnov et al. 1982
Sandlin et al. 1977
Feldman et al. 1998a
Peacock et al. 1984
Curdt et al. 2000
Peacock et al. 1984
Feldman et al. 1998a
Widing 1978
Kucera et al. 2000
Feldman et al. 1985
Feldman et al. 1998a
Feldman et al. 1998a
Feldman et al. 1998a
References to Previous
Measurements
5
5
5
3
6
7
4
5
6
Cross Reference
to Other Tables
1410.584 ^ 0.025
1432.228 ^ 0.050
1435.705 ^ 0.040/2 \ (717.853)
1443.107 ^ 0.025/2 \ (721.554)
1452.655 ^ 0.040
1460.877 ^ 0.030/2 \ (730.439)
1472.980 ^ 0.040/2 \ (736.490)
1481.479 ^ 0.020/2 \ (740.740)
1491.605 ^ 0.020/2 \ (745.803)
1504.275 ^ 0.032
1515.190 ^ 0.060
1517.180 ^ 0.040/2 \ (758.590)
1558.901 ^ 0.030/2 \ (779.451)
1572.061 ^ 0.030/2 \ (786.031)
1586.309 ^ 0.028 bl with
1586.309/2 \ (793.155)
1589.256 ^ 0.035/2 \ (794.628)
1392.095 ^ 0.017
1328.791 ^ 0.030
1330.532 ^ 0.030
1331.503 ^ 0.021
1340.603 ^ 0.020/2 \ (670.302)
1354.064 ^ 0.020
1354.892 ^ 0.025
1358.537 ^ 0.029/2 \ (679.269)
1375.959 ^ 0.035
730.41
736.55
740.75 ^ 0.03
717.88
721.55
1410.62 ^ 0.02
1392.12
679.29
1375.98
1330.54^0.02
1331.48^0.02
670.34^0.02
1354.1 ^ 0.1
Previously
Measured j
(AŽ )
6.45
6.45
6.2
6.9
6.2
6.2
6.2
6.6
6.2
6.45
6.2
6.9
6.9
6.95
6.9
6.75
B6.6
6.2
6.75
6.8
6.4
[6.2
6.3
6.95
6.75
6.9
6.55
log T b
e
(K)
Kink 1999
Kink 1999
Peacock et al. 1984
Kink 1999
Kucera et al. 2000
Kucera et al. 2000
Feldman & Doschek 1977
Sandlin et al. 1977
Feldman & Doschek 1977
Sandlin et al. 1977
Peacock et al. 1984
Feldman et al. 1998a
Feldman et al. 1997
Feldman et al. 1998a
Doschek et al. 1975
References to Previous
Measurements
2s22p3 2D È2s22p3 4S
3@2
3@2
2s22p4 3P È2s22p4 3P
779.5 ^ 0.3
Hinnov et al. 1982
1
2
2s22p2 1D È2s22p2 3P
786.1 ^ 0.3
Hinnov et al. 1982
2
2
2s22p3 2P È2s22p3 2P
1586.29^0.03
Kucera et al. 2000
3@2
1@2
2s22p3 2D È2s22p3 4S
793.3 ^ 0.3
Hinnov et al. 1982
3@2
3@2
2s22p53s 3P È2s22p53s 3P
0
1
a UnidentiÐed lines are noted by question marks or by temperature classiÐcation of Feldman et al. 1997, if available. Second-order lines are marked by /2.
b In all other than the He-like ions the temperature of maximum fractional abundance is given.
c For all He-like lines the given values represent the temperature at which the contribution function reaches maximum.
2s22p33p 5P È2s22p33s 5S
2
2
2s22p33p 5P È2s22p33s 5S
1
2
2s22p4 1D È2s22p4 3P
2
2
2s22p4 1S È2s22p4 3P
0
1
2s22p3 2P È2s22p3 2D
1@2
3@2
2s22p5 2P È2s22p5 3P
1@2
3@2
2s22p3 2P È2s22p3 2D
3@2
5@2
2s22p33p 5P È2s22p33s 5S
3
2
2s22p3 2D È2s22p3 4S
3@2
3@2
Cr XVI . . . . . . . .
Ca XIV . . . . . . . .
Ar XI . . . . . . . . . .
Fe XX . . . . . . . . .
? ..............
Ar XI . . . . . . . . . .
Ar XI . . . . . . . . . .
Cr XVII . . . . . . .
Ar XI . . . . . . . . . .
Ca XIV . . . . . . . .
? ..............
Mn XIX . . . . . . .
Ni XXI . . . . . . . .
Fe XXI . . . . . . . .
Fe XX] . . . . . .
Cr XVIII . . . . . . .
Ni XIX . . . . . . . .
2P È2s22p3 4S
1@2
3@2
3P È2s22p2 3P
1
0
3P È2s22p4 3P
1
2
2P È2s22p3 2D
3@2
5@2
1D È2s22p2 3P
2
2
2s22p4 1D È2s22p4 3P
2
2
2s22p3
2s22p2
2s22p4
2s22p3
2s22p2
2s22p4 1D È2s22p4 3P
2
2
2s22p4 3P È2s22p4 3P
0
2
2s22p2 1D È2s22p2 3P
2
1
Transition
Ar XI . . . . . . . . . .
Mn XVIII . . . . . .
Fe XIX . . . . . . . .
Ar XIII . . . . . . . .
(f) . . . . . . . . . . . . . .
Ar XII . . . . . . . . .
Fe XXI . . . . . . . .
Mn XVIII . . . . . .
Fe XX . . . . . . . . .
Ca XV . . . . . . . . .
Linea
Flare j
(AŽ )
TABLE 1ÈContinued
5
4
5
5
4
6
Cross Reference
to Other Tables
SPECTRAL LINES AND THEIR POTENTIAL USE IN PLASMA DIAGNOSTICS
TABLE 2
lengths and references to the quoted wavelengths. The temperatures at which the ions emitting the Ñare lines attain
their highest fractional abundance under coronal equilibrium conditions are also listed. Our wavelength measurements are given with uncertainty estimates.
In the following we discuss the high temperature lines
that are listed in Table 1. The discussion is arranged according to isoelectronic sequences.
4.1. He-like L ines
Transitions between levels belonging to the He-like 1s2p
3P and 1s2s 3S terms in all ions from Ne8` to Ca19` fall in
1
the 500È1600 AŽ range. Therefore, in principle, all lines originating from the transitions 1s2p 3P ] 1s2s 3S , 1s2p
3P ] 1s2s 3S , and 1s2p 3P ] 1s2s 3S2 should be1detect1
1
0
1
able in the SUMER range. However, because of the much
larger spontaneous decay rate of 1s2p 3P to the 1s2 1S
1
0
ground level than to 1s2p 3S , it is expected that the emis1
sivities of the 1s2p 3P ] 1s2s 3S lines is very low. They are
1 solar Ñare
1 spectra.
indeed not visible in the
The 1999 May 9 spectra contain the 1s2p 3P ] 1s2s 3S
2 The Ñare
1
transitions of Ne IX, Na X, Mg XI, and Si XIII.
spectra also include the much fainter 1s2p 3P ] 1s2s 3S
1
transitions of Ne IX, Mg XI, and Si XIII. Lines0 from other
fairly abundant elements S, Ar, and Ca that also fall in the
SUMER range have not been observed. The reason for not
detecting these lines is most likely due to the fact that the
temperature that existed in the plasma during the observations was not high enough to populate substantially the
1s2p 3P levels in S, Ar, and Ca.
Accurate wavelength measurements of the 1s2p
3P ] 1s2s 3S transitions are used as benchmarks for
0,2
1 calculations of relativistic and quantum
comparisons
with
electrodynamic (QED) e†ects in many body systems (Drake
1988 ; Plante, Johnson, & Sapirstein 1994). 1s2p
3P ] 1s2s 3S He-like transitions of Ne IX, Mg XI, Al XII,
0,2Si XIII were1 measured in the laboratory by beam-foil
and
techniques (e.g., Beyer, Folkmann, & Schartner 1986 ;
Brown et al. 1985 ; Howie et al. 1994 ; Klein et al. 1985). We
believe that our measurements of these wavelengths are the
Ðrst to be done from collisional astrophysical plasmas. The
actual wavelength determinations using the 1999 May 9
Ñare spectra and other solar spectra were described in more
detail by Curdt et al. (2000).
4.2. Be-like L ines
The Be-like spin forbidden transitions of the type 2s2p
3P ] 2s2 1S of solar abundant elements with Z º 14 are
1,2
0 in spectra emitted by quiet and active
commonly
seen
region plasmas (for details on such spectra see Curdt et al.
1997 ; Feldman et al. 1997). Be-like transitions of the type
2s2p 3P ] 2s2p 3P in ions with 24 ¹ Z ¹ 30 are also
expected2 to appear1 in the SUMER wavelength range.
Among the elements in the 24 ¹ Z ¹ 30 range, only Fe and
possibly Ni have solar abundances high enough to justify
expectations that their lines might be visible in the Ñare
513
ENERGY LEVELS IN Be-LIKE Fe XXIII
ConÐguration
Term
J
Level (cm~1)
2s2 . . . . . . . . . . . .
2s2p . . . . . . . . . . .
2s2p . . . . . . . . . . .
2s2p . . . . . . . . . . .
1S
3P
3P
3P
0
0
1
2
0
348238a
379125b
471768
a Predicted by Edlen 1983.
b Sugar & Rowan 1995.
spectra. In the 1999 May 9 Ñare only the 1079.415 AŽ Fe XXIII
line was detected presumably because the plasma temperature was not sufficiently high for a signiÐcant population of Be-like Ni ions to exist. This line was not seen
before in solar plasmas. Hinnov et al. (1982) detected this
line in tokamak spectra at 1079.3 AŽ using a spectrometer
with modest resolution. The energy level derived from our
measurement is given in Table 2.
4.3. B-like L ines
The SUMER wavelength range contains a large number
of B-like lines associated with transitions of the type
2s2p2 ] 2s22p from elements with Z ¹ 16. Many of these
lines are very prominent in quiet region, coronal hole, and
active region spectra (see Curdt et al. 1997 ; Feldman et al.
1997). The only bright lines from B-like ions that are associated with high-temperature plasmas (T º 2 ] 106 K) and
e
are expected to appear in the SUMER
spectral range
belong to the 2s22p 2P ] 2s22p 2P
forbidden tran1@2Fe XXII forbidden
sition in elements with 243@2¹ Z ¹ 29. The
transition was Ðrst observed in solar Ñare spectra obtained
with the OSO 6 spectrometer (Doschek et al. 1975). Hinnov
& Suckewer (1980) also detected the same line in tokamak
spectra with a wavelength of 845.55 ^ 0.1 AŽ . Hinnov et al.
(1982) measured a line at 1205.9 ^ 0.3 AŽ in tokamak spectra
and a second line at 609.9 ^ 0.3 AŽ , which they identiÐed as
the 2s22p 2P ] 2s22p 2P forbidden transition in Cr XX
and Ni XXIV,3@2
respectively. 1@2
We have measured in our spectra the very bright Fe XXII
2s22p 2P ] 2s2 2P
forbidden line at 845.570 AŽ . In
addition, 3@2
we detected1@2the much fainter Cr XX line at
1205.728 AŽ and newly identiÐed the Mn XXI line at 1005.862
AŽ . The Ni XXIV line is predicted to appear in close wavelength proximity to the very bright Mg X 609 AŽ resonance
line. It is likely that because of the overwhelming brightness
of the Mg X line, the Ni XXIV line could not be resolved from
it with certainty. The energy levels within the ground conÐguration of Cr XX, Mn XXI, and Fe XXII resulting from our
measurements are given in Table 3.
4.4. C-like L ines
The SUMER spectral range contains lines that originate
in transitions between the Ðrst excited conÐguration 2s2p3
TABLE 3
2s22p ENERGY LEVELS IN B-LIKE Cr XX, Mn XXI, AND Fe XXII
ConÐguration
Term
J
Cr XX level (cm~1)
Mn XXI level (cm~1)
Fe XXII level (cm~1)
2s22p . . . . . . . . . .
2s22p . . . . . . . . . .
2P
2P
1/2
3/2
0
82937.4 ^ 1.7
0
99417.2 ^ 2.0
0
118263.4 ^ 2.8
514
FELDMAN ET AL.
Vol. 544
TABLE 4
TABLE 6
2s22p2 ENERGY LEVELS IN C-LIKE Fe XXI
2s22p4 ENERGY LEVELS IN O-like Fe XIX
ConÐguration
2s22p2
2s22p2
2s22p2
2s22p2
2s22p2
Term
J
Level (cm~1)
3P
3P
3P
1D
1S
0
1
2
2
0
0
73851.8 ^ 1.1
117347.0 ^ 6.0
244568.4 ^ 5.0
371954a
........
........
........
........
........
ConÐguration
2s22p4
2s22p4
2s22p4
2s22p4
2s22p4
a Predicted by Edlen 1985.
........
........
........
........
........
Term
J
Level (cm~1)
3P
3P
3P
1D
1S
2
0
1
2
0
0
75256.4 ^ 1.7
89439.2 ^ 4.3
168852.3 ^ 4.3
325140 ^ 5a
a Widing 1978.
and the ground conÐguration 2s22p2 in elements with
Z \ 15. In addition our range contains forbidden lines of
Ar XIII, Ca XV, Cr XIX, Fe XXI, and Ni XXIII that arise from
transition within the 2s22p2 ground conÐguration. All transitions that were detected in the 1999 May 9 Ñare with the
exception of Ca XV 555.382 AŽ have already been seen before
either in solar spectra (e.g., Doschek et al. 1975 ; Feldman et
al. 1998a) or in tokamak plasma (Hinnov et al. 1982).
However, we signiÐcantly improved on the wavelength
determinations of two of the Fe XXI lines and the Ni XXIII
line. The resulting energy levels of Fe XXI are given in
Table 4.
4.5. N-like L ines
We expect in our spectra bright lines from elements with
Z \ 18 that originate from transitions between the 2s22p3
ground conÐguration and the Ðrst excited conÐguration
2s2p4. The lines, which originate from such transitions, are
emitted by relatively low-temperature astrophysical
plasmas. N-like transition that originate within the 2s22p3
ground conÐguration in elements with 18 º Z º 30 also
appear in high-temperature SUMER spectra. The Ðve levels
in the N-like ground conÐguration produce nine di†erent
forbidden transitions. In an earlier publication (Feldman et
al. 1998a) we presented identiÐcations of Ar XII and Ca XIV
lines that were recorded from a hot active region. In a
second publication Kucera et al. (2000) presented a comprehensive list of Fe XX forbidden lines that were based on
solar Ñares and on tokamak spectra.
Because of the very high brightness of the 1999 May 9
Ñare, we also succeeded in measuring the previously undetected 649.093 AŽ Ar XII line and the 1432.228 and 1504.275
AŽ Ca XIV lines. We also measured the two previously
unidentiÐed K XIII lines and a Co XXI line also belonging to
the N I isoelectronic sequence. In addition, we improved the
wavelengths of Ti XVI and Cr XVIII lines previously measured in tokamak spectra. Energy levels within the ground
conÐguration in Ðve di†erent ions are listed in Table 5.
4.6. O-like L ines
The SUMER Ñare spectra contain a number of forbidden
lines that arise from transitions within the ground conÐguration 2s22p4 in O-like Ar, Ca, Ti, Cr, Mn, Fe, and Ni ions.
Four of the lines, the Ar XI line at 745.803 AŽ , the two Mn
XVIII lines at 1354.892 and 662.951 AŽ , and the Fe XIX line at
1328.791 AŽ , were observed for the Ðrst time in the Ñare
spectra. The other lines have already been seen in solar (e.g.,
Feldman & Doschek 1977 ; Sandlin et al. 1977 ; Feldman et
al. 1998a) and in tokamak spectra (Hinnov et al. 1982 ;
Peacock et al. 1984). We used the newly measured Fe XIX
line at 1328.791 AŽ to improve the energy levels of the
ground conÐguration given in Table 6.
4.7. F-like L ines
The SUMER wavelength range contains the 2s22p5
2P ] 2s22p5 2P
transition from elements with
3@2 et al. (1975) Ðrst reported the pres24 1@2
¹ Z ¹ 30. Doschek
ence of the Fe XVIII line in solar coronal spectra. Subsequently, measurements of the Fe XVIII, together with Cr XVI
and Ni XX, were reported in tokamak spectra. Peacock et al.
(1984) claimed the most accurate measurement so far of the
2s22p5 2P ] 2s22p5 2P transition in Cr XVI 1410.62 AŽ ,
1@2 AŽ , and in 3@2
Fe XVIII 974.86
Ni XX 694.64 AŽ . We have detected
2s22p5 2P ] 2s22p5 2P
transition in Cr XVI and Fe
3@2 We could not detect the Ni XX
XVIII in the1@2
1999 May 9 Ñare.
transition possibly because of its blend with a bright Si IX
line. However, we detected for the Ðrst time the 2s22p5
2P ] 2s22p5 2P
transition in Mn XVII at 1167.760 AŽ
1@2in Co XIX at 820.019
3@2
and
AŽ . The newly determined energy
levels of Mn XVII and Co XIX are given in Table 7.
4.8. Ne-like L ines
Ne-like resonance transitions between the ground conÐguration 2s22p6 1S and the Ðrst excited conÐguration
0
2s22p53s in all ions other
than Ne I have wavelengths too
short to fall in the SUMER spectral range. The 2s22p53s
conÐguration consists of the four levels, 3P , 3P , 3P , and
2
1
0
TABLE 5
2s22p3 ENERGY LEVELS IN N-like Ar XII, K XIII, Ca XIV, Ti XVI, AND Cr XVIII
ConÐguration
2s22p3
2s22p3
2s22p3
2s22p3
2s22p3
........
........
........
........
........
Term
J
Ar XII Level (cm~1)
K XIII Level (cm~1)
Ca XIV Level (cm~1)
Ti XVI Level (cm~1)
Cr XVIII Level (cm~1)
4S
2D
2D
2P
2P
3/2
3/2
5/2
1/2
3/2
0
94824.0 ^ 1.3
98155.7 ^ 1.9
149186.5 ^ 2.2
154061.1 ^ 4.7
0
100558.9 ^ 2.0
105721.7 ^ 2.2
160640b
168084b
0
105978.8 ^ 2.2
113584.6 ^ 1.9
172457.5 ^ 4.5
183408.0 ^ 5.0
0
116036.3 ^ 2.7
130701a
197668c
219249.5 ^ 4.0
0
126078.0 ^ 3.3
150840.7 ^ 3.4
226193d
264550d
a Predicted by Edlen 1984.
b Predicted by Edlen 1985.
c E. Hinnov, unpublished 1986.
d Denne & Hinnov 1984.
No. 1, 2000
SPECTRAL LINES AND THEIR POTENTIAL USE IN PLASMA DIAGNOSTICS
515
TABLE 7
2s22p5 ENERGY LEVELS IN F-like Mn XVII AND Co XIX
ConÐguration
Term
J
Mn XVII Level (cm~1)
Co XIX Level (cm~1)
2s22p5 . . . . . . . .
2s22p5 . . . . . . . .
2P
2P
3/2
1/2
0
85634.0 ^ 1.1
0
121948.3 ^ 3.7
1P , where 3P has the lowest energy and 1P the highest.
1
2
1
At solar densities the main depopulation mechanism of 3P ,
2
3P , and 1P is via radiative decay into the ground level.
1
1
Since a direct radiative decay from the 3P level to the 1S
0 primarily as 0a
ground level is strictly forbidden, it decays
forbidden (M1) transition to the 3P level at solar electron
1
densities. Using Skylab Ñare spectra, Feldman, Doschek, &
Seely (1985) measured a line at 1135.20 AŽ which they identiÐed as the Fe XVII 2s22p53s 3P ] 2s22p53s 3P transition.
0
1
Using the 1999 May 9 Ñare we also measured a hightemperature line at 794.605 AŽ , which we identiÐed as the Ni
XIX 2s22p53s 3P ] 2s22p53s 3P transition (see Table 1).
0
1
This newly identiÐed
line establishes
the energy separation
between the 3P and 3P levels as 125849 cm~1.
1
0
5.
HIGH-TEMPERATURE, DENSITY-SENSITIVE LINE RATIOS
The SUMER spectral range contains a number of line
ratios suitable for determining the density of hightemperature solar plasmas. A list of the line ratios, the electron density range over which each intensity ratio is
sensitive and the temperature at which the ions emitting the
lines reach their maximum fractional abundances is given in
Table 8. Figure 2 displays line ratios as a function of electron density for several of the Ar XII, K XIII, and Ca XIV line
pairs. (The intensity ratios are displayed in energy units and
not in photon units.) As seen from the Ðgure the intensity
ratios, typically emitted by 2 ] 106 to 4 ] 106 K plasmas,
could be used to derive electron densities in the 1 ] 109 to
1 ] 1013 cm~3 ranges. These ratios have been calculated
using the data of the CHIANTI database (Dere et al. 1997 ;
Landi et al. 1999).
High-temperature Fe line ratios, of which the Fe XXI
shown in Figure 3 is a typical example, are sensitive to
plasmas with densities higher than 1 ] 1011 cm~3. Important exceptions to the rule are Fe XIX line ratios involving
the 2s22p4 3P ] 2s22p4 3P transition. As seen from
2 are sensitive to densities of
Figure 4, the Fe0 XIX line ratios
1 ] 109 cm~1 and higher.
6.
TABLE 8
HIGH-TEMPERATURE DENSITY-SENSITIVE LINE RATIOS IN THE
SUMER WAVELENGTH RANGE
Line 1
(AŽ )
Line 2
(AŽ )
log density
(cm~3)
XII . . . . . . .
XII . . . . . . .
XII . . . . . . .
XII . . . . . . .
XII . . . . . . .
1018
1054
1054
1018
1018
1054
670
649
670
649
9È11.5
9È13
10È12.5
9È12
8.5È12.5
log T (K)a
e
6.3
6.3
6.3
6.3
6.3
Ar XIII . . . . . .
1330
656
[10.5
6.4
K XIII . . . . . . .
945
994
9È12
6.35
Ca XIII . . . . . .
1133
648
[11
6.35
XIV . . . . . .
XIV . . . . . .
XIV . . . . . .
XIV . . . . . .
XIV . . . . . .
XIV . . . . . .
XIV . . . . . .
XIV . . . . . .
545
545
880
580
580
880
880
943
880
944
580
944
1291
944
1291
1291
10È13
10È13
10È13
10È13
10È13
10È12
10È13
10È13
6.45
6.45
6.45
6.45
6.45
6.45
6.45
6.45
Ca XV . . . . . .
Ca XV . . . . . .
1098
1375
555
555
[11
[11
6.55
6.55
Fe XIX . . . . . .
Fe XIX . . . . . .
1329
1329
1118
592
[9
[9
6.8
6.8
Fe XXI . . . . . .
786
1354
[11
6.95
Line
Ar
Ar
Ar
Ar
Ar
Ca
Ca
Ca
Ca
Ca
Ca
Ca
Ca
a Temperature of maximum fractional abundance of the ion.
EMISSIVITIES OF HIGH-TEMPERATURE LINES
The SUMER wavelength range includes lines that could
be used to determine emission measure distributions as well
as the plasma temperatures in the 2 ] 106È2 ] 107 K range.
In particular, lines of Ca XIII, Ca XIV, Ca XV, which have
their maximum fractional abundances at log T of 6.35,
e Fe XX,
6.45, 6.55, respectively, and lines of Fe XVIII, Fe XIX,
Fe XXII, Fe XXIII, which have their maximum fractional
abundance at log T /k of 6.7, 6.8, 6.9, 6.95, 7.05, 7.1 (Arnaud
e respectively, can be used for emission
& RothenÑug 1985),
measure and plasma temperature determinations. Table 9
provides emissivities of Ca lines in the log T \ 6.2È6.8
e
range and log N in the 9È13 range. Table 10 provides
emise
sivities of Fe lines in the log T \ 6.6È7.3 range and log N
e calculations given in thee
in the 9È13 range. The emissivity
tables are based on cross sections stored in the CHIANTI
FIG. 2.ÈIntensity ratios (in energy units) vs. electron density for Ar XII,
K XIII, and Ca XIV lines.
FIG. 3.ÈIntensity ratios (in energy units) vs. electron density for the
786/1354 Fe XXI lines.
FIG. 4.ÈIntensity ratios (in energy units) vs. electron density for the
1328/1118 Fe XIX lines.
TABLE 9
EMISSIVITIES FOR Ca XIII, Ca XIV, AND Ca XV LINES APPEARING IN 500È1600 AŽ RANGE
LINE
log EMISSIVITY1 (ergs cm~3 s~1) VERSUS log T (K)
e
j
(AŽ )
log N
e
(cm~3)
6.2
6.3
6.4
6.5
6.6
6.7
6.8
Ca
Ca
Ca
Ca
Ca
XIII . . . . . .
XIII . . . . . .
XIII . . . . . .
XIII . . . . . .
XIII . . . . . .
1133
1133
1133
1133
1133
9
10
11
12
13
15.810
15.812
15.836
16.025
16.663
15.219
15.222
15.244
15.422
16.037
14.962
14.965
14.986
15.152
15.743
15.136
15.140
15.160
15.314
15.880
15.692
15.695
15.715
15.858
16.396
16.589
16.594
16.613
16.745
17.256
17.841
17.846
17.865
17.987
18.470
Ca
Ca
Ca
Ca
Ca
XIII . . . . . .
XIII . . . . . .
XIII . . . . . .
XIII . . . . . .
XIII . . . . . .
648
648
648
648
648
9
10
11
12
13
16.502
16.487
16.438
16.307
16.457
15.511
15.896
15.848
15.710
15.835
15.655
15.640
15.592
15.448
15.547
15.832
15.817
15.770
15.621
15.692
16.392
16.376
16.331
16.178
16.219
17.295
17.279
17.236
17.079
17.090
18.544
18.358
18.495
18.336
18.316
Ca
Ca
Ca
Ca
Ca
XIV . . . . . .
XIV . . . . . .
XIV . . . . . .
XIV . . . . . .
XIV . . . . . .
943
943
943
943
943
9
10
11
12
13
16.721
16.731
16.805
17.193
18.016
15.551
15.559
15.624
15.992
16.795
14.817
14.822
14.879
15.225
16.005
14.595
14.598
14.646
14.969
15.724
14.814
14.816
14.856
15.156
15.884
15.423
15.423
15.455
15.733
16.432
16.423
16.421
16.446
16.701
17.371
Ca
Ca
Ca
Ca
Ca
XIV . . . . . .
XIV . . . . . .
XIV . . . . . .
XIV . . . . . .
XIV . . . . . .
880
880
880
880
880
9
10
11
12
13
16.922
17.071
17.646
18.491
19.389
15.749
15.890
16.448
17.282
18.166
15.012
15.143
15.684
16.507
17.376
14.786
14.908
15.430
16.242
17.094
15.002
15.115
15.617
16.418
17.253
15.608
15.712
16.193
16.983
17.800
16.604
16.699
17.159
17.938
18.738
Ca
Ca
Ca
Ca
Ca
XIV . . . . . .
XIV . . . . . .
XIV . . . . . .
XIV . . . . . .
XIV . . . . . .
579
579
579
579
579
9
10
11
12
13
17.314
17.286
17.222
17.266
17.828
16.144
16.115
16.044
16.069
16.608
15.411
15.381
15.304
15.309
15.820
15.190
15.159
15.077
15.062
15.541
15.413
15.382
15.294
15.259
15.704
16.025
15.994
15.901
15.847
16.256
17.030
16.999
16.901
16.829
17.199
Ca
Ca
Ca
Ca
Ca
XIV . . . . . .
XIV . . . . . .
XIV . . . . . .
XIV . . . . . .
XIV . . . . . .
1504
1504
1504
1504
1504
9
10
11
12
13
18.197
18.169
18.105
18.149
18.711
17.027
16.998
16.927
16.952
17.491
16.294
16.264
16.187
16.192
16.703
16.073
16.043
15.960
15.945
16.424
16.296
16.265
16.177
16.142
16.587
16.909
16.877
16.784
16.730
17.139
17.913
17.882
17.784
17.712
18.082
Ca
Ca
Ca
Ca
Ca
XIV . . . . . .
XIV . . . . . .
XIV . . . . . .
XIV . . . . . .
XIV . . . . . .
545
545
545
545
545
9
10
11
12
13
17.138
17.067
16.946
16.910
17.306
15.965
15.892
15.762
15.713
16.086
15.227
15.154
15.016
14.952
15.299
15.003
14.929
14.784
14.706
15.023
15.221
15.147
14.995
14.903
15.188
15.829
15.756
15.598
15.493
15.743
16.828
16.757
16.592
16.474
16.690
SPECTRAL LINES AND THEIR POTENTIAL USE IN PLASMA DIAGNOSTICS
517
TABLE 9ÈContinued
LINE
log EMISSIVITY1 (ergs cm~3 s~1) VERSUS log T (K)
e
j
(AŽ )
log N
e
(cm~3)
6.2
6.3
6.4
6.5
6.6
6.7
6.8
Ca
Ca
Ca
Ca
Ca
XIV . . . . . .
XIV . . . . . .
XIV . . . . . .
XIV . . . . . .
XIV . . . . . .
1291
1291
1291
1291
1291
9
10
11
12
13
17.846
17.775
17.654
17.618
18.014
16.673
16.600
16.470
16.421
16.794
15.935
15.862
15.724
15.660
16.007
15.711
15.637
15.492
15.414
15.731
15.929
15.855
15.703
15.611
15.896
16.537
16.464
16.306
16.201
16.451
17.536
17.465
17.300
17.182
17.398
Ca
Ca
Ca
Ca
Ca
XIV . . . . . .
XIV . . . . . .
XIV . . . . . .
XIV . . . . . .
XIV . . . . . .
1432
1432
1432
1432
1432
9
10
11
12
13
18.264
18.192
18.071
18.035
18.432
17.090
17.017
16.887
16.838
17.212
16.353
16.279
16.141
16.078
16.425
16.128
16.054
15.909
15.831
16.148
16.346
16.273
16.121
16.029
16.313
16.954
16.882
16.723
16.618
16.868
17.954
17.882
17.717
17.600
17.815
Ca
Ca
Ca
Ca
Ca
XV
XV
XV
XV
XV
......
......
......
......
......
1098
1098
1098
1098
1098
9
10
11
12
13
17.492
17.473
17.438
17.571
18.130
16.218
16.195
16.147
16.265
16.803
15.537
15.510
15.448
15.552
16.071
15.361
15.330
15.253
15.343
15.846
15.621
15.585
15.493
15.570
16.058
16.310
16.269
16.162
16.226
16.698
Ca
Ca
Ca
Ca
Ca
XV
XV
XV
XV
XV
......
......
......
......
......
1375
1375
1375
1375
1375
9
10
11
12
13
17.420
17.401
17.366
17.499
18.058
16.146
16.123
16.075
16.193
16.731
15.465
15.438
15.376
15.480
15.999
15.289
15.258
15.181
15.271
15.774
15.549
15.513
15.421
15.498
15.986
16.238
16.197
16.090
16.154
16.626
Ca
Ca
Ca
Ca
Ca
XV
XV
XV
XV
XV
......
......
......
......
......
555
555
555
555
555
9
10
11
12
13
17.680
17.671
17.630
17.493
17.619
16.383
16.374
16.334
16.195
16.300
15.678
15.670
15.632
15.491
15.576
15.476
15.469
15.436
15.294
15.358
15.707
15.703
15.677
15.533
15.578
16.367
16.365
16.346
16.203
16.228
a Data for the CHIANTI calculations was taken from the following sources : Ca XIII : Radiative data from
unpublished SUPERSTRUCTURE calculations by A. K. Bhatia 1996, private communication ; collisional data
taken from unpublished distorted wave calculations by A. K. Bhatia 1996, private communication. Ca XIV :
Radiative and collisional data from Bhatia & Mason 1980a. Ca XV : Radiation data from Froese Fischer & Saha
1985 ; collisional data from Aggarwal, Berrington & Keenan 1991 ; Bhatia & Doschek 1993.
database, on the fractional ionization calculations of Mazzotta et al. (1998), and on photospheric abundances given in
Grevesse & Sauval (1998).
7.
TEMPERATURE DIAGNOSTICS
Plasma temperatures could be determined using the lines
in the SUMER wavelength range. Assuming that the Ñaring
plasma is in a steady state coronal equilibrium, the temperature of the plasma could be derived as follows :
In one approach the determination is based on intensity
ratios of lines emitted by the same ion. At Ðrst, the electron
density is determined using the line ratios described in ° 4.
Given the electron density, the plasma temperature could
be determined by comparing measured intensity ratios of
lines belonging to the same ion with calculations at a
number of temperatures. In particular, the intensity ratios
of the Fe XIX 1328/592 and 1118/592 and the intensity ratios
of the Ca XV 1098/555 and 1375/555 could be used for the
purpose.
In a second method calculated line emissivities versus
temperature are divided by the measured line intensities.
Using the derived values, curves representing the ratios as a
function of temperature are generated for the measured
lines. (Emissivities as a function of temperature and density
for Ca XIIIÈCa XV and Fe XVIIIÈFe XXIII are listed in Tables 9
and 10.) In cases where the plasma under consideration is
isothermal, the curves from each of the lines, independent of
their degree of ionization, should cross each other in nearly
the same point. The point of intersection could be considered as the plasma temperature. For illustration of this
approach, see Figure 4 in Feldman et al. (1999).
A third method for determining the plasma temperature
involves a comparison of intensity ratios of lines originating
from very highly excited levels with lines originating in lowexcitation levels. In particular, this technique could work
well in comparing the intensity of the 1153 AŽ Fe XVII line
that originates in a level which is 725 eV above the ground
state with Fe XVIIIÈFe XXII lines originating from levels in
the ground conÐguration. A second interesting system
involves a comparison of the 1248 AŽ Ne IX line which originates in a level that is 915 eV above the ground state, with
Ar XIIÈAr XIV lines originating in the ground conÐguration.
Since both Ne and Ar are high Ðrst ionization potential
(FIP) elements (see ° 8) their abundances are expected, anywhere in the solar upper atmosphere, to be photospheric.
TABLE 10
EMISSIVITIES FOR FeE XVIIIÈFe XXXI LINES APPEARING IN 500È1600 AŽ RANGE
LINE
log EMISSIVITY1 (ergs cm~3 s~1) VERSUS log T (K)
e
j
(AŽ )
log N
e
(cm~3)
6.6
6.7
6.8
6.9
7.0
7.1
7.2
7.3
974
974
974
974
974
9
10
11
12
13
14.259
14.259
14.260
14.266
14.322
13.926
13.926
13.927
13.933
13.991
13.746
13.746
13.747
13.753
13.812
13.813
13.813
13.813
13.820
13.880
14.256
14.256
14.257
14.264
14.325
15.125
15.125
15.125
15.132
15.193
16.341
16.341
16.342
16.348
16.409
17.690
17.690
17.691
17.697
17.757
7.4
Fe
Fe
Fe
Fe
Fe
XVIII . . . . . .
XVIII . . . . . .
XVIII . . . . . .
XVIII . . . . . .
XVIII . . . . . .
Fe
Fe
Fe
Fe
Fe
XIX
XIX
XIX
XIX
XIX
.......
.......
.......
.......
.......
1328
1328
1328
1328
1328
9
10
11
12
13
16.216
16.986
17.954
18.951
19.956
15.576
16.339
17.306
18.303
19.307
15.150
15.905
16.871
17.867
18.870
15.009
15.755
16.719
17.716
18.718
15.270
16.007
16.970
17.966
18.967
15.972
16.699
17.659
18.655
19.655
17.029
17.745
18.703
19.698
20.698
18.225
18.928
19.884
20.879
21.877
19.407
20.098
21.051
22.045
23.043
Fe
Fe
Fe
Fe
Fe
XIX
XIX
XIX
XIX
XIX
.......
.......
.......
.......
.......
1118
1118
1118
1118
1118
9
10
11
12
13
15.323
15.318
15.318
15.334
15.468
14.694
14.688
14.688
14.703
14.832
14.277
14.270
14.270
14.284
14.408
14.145
14.136
14.136
14.150
14.268
14.415
14.405
14.405
14.417
14.530
15.124
15.114
15.113
15.125
15.233
16.190
16.179
16.178
16.189
16.292
17.394
17.383
17.381
17.392
17.489
18.586
18.573
18.572
18.582
18.674
Fe
Fe
Fe
Fe
Fe
XIX
XIX
XIX
XIX
XIX
.......
.......
.......
.......
.......
592
592
592
592
592
9
10
11
12
13
15.354
15.355
15.356
15.363
15.426
14.730
14.731
14.731
14.738
14.798
14.316
14.316
14.317
14.323
14.380
14.185
14.185
14.186
14.192
14.245
14.458
14.457
14.458
14.463
14.513
15.170
15.170
15.170
15.175
15.222
16.241
16.240
16.240
16.245
16.288
17.450
17.449
17.449
17.454
17.493
18.647
18.646
18.646
18.650
18.687
Fe
Fe
Fe
Fe
Fe
XX . . . . . . . .
XX . . . . . . . .
XX . . . . . . . .
XX . . . . . . . .
XX . . . . . . . .
721
721
721
721
721
9
10
11
12
13
16.520
16.520
16.521
16.535
16.641
15.514
15.514
15.516
15.528
15.627
14.780
14.780
14.781
14.792
14.884
14.369
14.369
14.370
14.380
14.466
14.387
14.387
14.388
14.397
14.478
14.866
14.866
14.867
14.876
14.952
15.719
15.719
15.720
15.728
15.799
16.727
16.727
16.728
16.735
16.802
17.734
17.734
17.735
17.741
17.804
Fe
Fe
Fe
Fe
Fe
XX . . . . . . . .
XX . . . . . . . .
XX . . . . . . . .
XX . . . . . . . .
XX . . . . . . . .
567
567
567
567
567
9
10
11
12
13
16.734
16.734
16.743
16.817
17.211
15.731
15.732
15.739
15.809
16.190
14.999
15.000
15.007
15.073
15.441
14.590
14.590
14.597
14.660
15.016
14.609
14.610
14.617
14.676
15.020
15.090
15.090
15.096
15.153
15.486
15.943
15.944
15.950
16.003
16.325
16.952
16.952
16.958
17.009
17.318
17.960
17.960
17.965
18.013
18.311
Fe
Fe
Fe
Fe
Fe
XX . . . . . . . .
XX . . . . . . . .
XX . . . . . . . .
XX . . . . . . . .
XX . . . . . . . .
821
821
821
821
821
9
10
11
12
13
18.111
18.110
18.108
18.088
18.032
17.105
17.105
17.103
17.082
17.021
16.369
16.369
16.366
16.346
16.283
15.954
15.954
15.952
15.932
15.867
15.968
15.967
15.965
15.947
15.882
16.442
16.442
16.440
16.422
16.358
17.290
17.290
17.288
17.272
17.209
18.294
18.294
18.292
18.276
18.215
19.297
19.297
19.295
19.281
19.220
Fe
Fe
Fe
Fe
Fe
XX . . . . . . . .
XX . . . . . . . .
XX . . . . . . . .
XX . . . . . . . .
XX . . . . . . . .
679
679
679
679
679
9
10
11
12
13
18.095
18.094
18.085
18.015
17.833
17.097
17.097
17.088
17.018
16.828
16.373
16.372
16.363
16.294
16.098
15.971
15.971
15.962
15.893
15.692
16.000
15.999
15.990
15.923
15.717
16.489
16.489
16.480
16.414
16.204
17.352
17.351
17.343
17.278
17.066
18.370
18.369
18.361
18.297
18.083
19.386
19.385
19.377
19.315
19.100
Fe
Fe
Fe
Fe
Fe
XX . . . . . . . .
XX . . . . . . . .
XX . . . . . . . .
XX . . . . . . . .
XX . . . . . . . .
1586
1586
1586
1586
1586
9
10
11
12
13
19.374
19.373
19.363
19.294
19.112
18.376
18.375
18.336
18.297
18.107
17.651
17.651
17.642
17.573
17.377
17.250
17.249
17.241
17.172
16.971
17.279
17.278
17.269
17.202
16.996
17.768
17.767
17.759
17.693
17.483
18.631
18.360
18.622
18.557
18.345
19.648
19.648
19.640
19.576
19.362
20.665
20.664
20.656
20.594
20.379
Fe
Fe
Fe
Fe
Fe
XXI
XXI
XXI
XXI
XXI
.......
.......
.......
.......
.......
1354
1354
1354
1354
1354
9
10
11
12
13
15.834
15.835
15.581
15.964
16.383
14.591
14.593
14.607
14.711
15.111
13.741
13.743
13.755
13.852
14.233
13.380
13.381
13.392
13.481
13.845
13.527
13.528
13.538
13.620
13.967
14.088
14.089
14.098
14.173
14.504
14.836
14.837
14.845
14.914
15.229
15.608
15.609
15.617
15.680
15.979
Fe XXI . . . . . . .
Fe XXI . . . . . . .
Fe XXI . . . . . . .
585
585
585
9
10
11
17.034
17.026
16.965
15.797
15.791
15.736
14.954
14.948
14.899
14.601
14.596
14.551
14.759
14.754
14.713
15.332
15.328
15.289
16.094
16.090
16.053
16.881
16.877
16.843
SPECTRAL LINES AND THEIR POTENTIAL USE IN PLASMA DIAGNOSTICS
519
TABLE 10ÈContinued
LINE
j
(AŽ )
log N
e
(cm~3)
Fe XXI . . . . . . .
Fe XXI . . . . . . .
585
585
log EMISSIVITY1 (ergs cm~3 s~1) VERSUS log T (K)
e
6.7
6.8
6.9
7.0
7.1
7.2
7.3
7.4
12
13
16.713
16.538
15.496
15.306
14.669
14.464
14.330
14.110
14.499
14.265
15.081
14.833
15.851
15.588
16.646
16.369
17.182
17.174
17.112
16.861
16.686
15.945
15.939
15.884
15.644
15.453
15.101
15.096
15.046
14.817
14.612
14.749
14.744
14.699
14.478
14.258
14.907
14.902
14.861
14.647
14.413
15.480
15.475
15.437
15.229
14.981
16.241
16.238
16.201
15.999
15.736
17.029
17.025
16.991
16.793
16.517
15.468
15.468
15.469
15.482
15.598
14.238
14.238
14.239
14.252
14.361
13.547
13.548
13.549
13.561
13.665
13.405
13.405
13.406
13.417
13.516
13.705
13.705
13.706
13.717
13.810
14.215
14.215
14.216
14.226
14.314
14.770
14.771
14.772
14.781
14.863
16.172
16.172
16.173
16.185
16.287
15.136
15.136
15.137
15.148
15.246
14.697
14.697
14.698
14.709
14.802
14.741
14.741
14.742
14.752
14.841
15.029
15.029
15.030
15.039
15.123
15.387
15.387
15.388
15.397
15.477
Fe
Fe
Fe
Fe
Fe
XXI
XXI
XXI
XXI
XXI
.......
.......
.......
.......
.......
786
786
786
786
786
9
10
11
12
13
Fe
Fe
Fe
Fe
Fe
XXII
XXII
XXII
XXII
XXII
......
......
......
......
......
845
845
845
845
845
9
10
11
12
13
Fe
Fe
Fe
Fe
Fe
XXIII . . . . . .
XXIII . . . . . .
XXIII . . . . . .
XXIII . . . . . .
XXIII . . . . . .
1079
1079
1079
1079
1079
9
10
11
12
13
6.6
a Data for the CHIANTI calculations was taken from the following sources :
Fe XVII : Radiative and collisional data from Bhatia & Doschek 1992.
Fe XVIII : Radiative data from Sampson et al. 1991 ; Blackford & Hibbert 1994 ; Cornille et al. 1992. Collisional data from Sampson et al. 1991.
Fe XIX : Radiative and collisional data from Lolergue et al. 1985 ; Bhatia et al. 1989a.
Fe XX : Radiative data from Bhatia et al. 1989b ; collisional data from B. McLaughlin 2000, private communication ; Bhatia & Mason 1980b.
Fe XXI : Radiative data from Young (1996, unpublished) ; Collisional data from Aggarwal 1991. Fe XXII : Radiative data from Dankwort &
Tre†tz 1978 ; Nussbaumer & Storey 1981 ; Lennon et al. 1985 ; Collisional data from Zhang, Graziani & Pradhan 1994. Fe XXIII : Radiative
and collisional data from Zhang & Sampson 1992.
LINES SUITABLE FOR DETERMINING ELEMENTAL
ABUNDANCES IN HIGH-TEMPERATURE PLASMAS
8.
During the last three decades instruments mounted on a
variety of spacecrafts gathered a large body of data regarding the properties of the solar wind. In analyzing the data it
was found that the composition of the solar wind is not
always identical to the composition of the solar photosphere. While the composition in some solar wind streams
(fast speed solar wind) resembles photospheric composition,
in other locations (in particular, in the slow speed solar
wind) it is modiÐed (e.g., von Steiger, Geiss, & Gloeckler
1997). The modiÐcation patterns were such that in general,
elements (excluding He) with FIP larger than 11.5 eV have
abundances similar to those in the photosphere. On the
other hand, the abundances of elements having FIP \ 10
eV in general are increased by factors of 4È5 relative to the
photospheric abundance. Elements having FIP [ 11.5 eV
were named high-FIP elements and those with FIP \ 10
eV were named low-FIP elements. The factor by which the
low-FIP elements are modiÐed relative to the photosphere
was deÐned as the FIP bias.
The literature contains well-documented examples of
plasmas with low-FIP/high-FIP composition ratios that
very between photospheric and 15 times photospheric.
Assuming that the modiÐcation occurs in the low-FIP elements while the high FIP elements stay unchanged, a question that is not yet completely settled, the FIP-bias can vary
between 1 and 15. A summary of the latest results regarding
elemental abundance modiÐcations in various solar regions
is presented in a number of review papers (e.g., Feldman
1992 ; Saba 1995 ; Meyer 1996 ; Feldman & Laming 2000).
Using SUMER spectra, Feldman et al. (1998b) derived
the abundance in the coronal plasma above a quiet region
(1.3 ] 106 K) and above a coronal hole region (9 ] 105 K).
They found that at heights of 30@@ above the limb in coronal
holes the composition of elements is very similar to the
composition of the solar photosphere. In contrast at 30@@
above the limb in quiet regions the low-FIP elements were
enhanced relative to their abundance in the photosphere by
a factor of 3È4. Dwivedi, Curdt, & Wilhelm (1999) found
that the FIP bias across an active region is increasing with
height, a trend that is confused by local excursions in the
active region loops. All earlier measurements were done by
X-ray instruments that were unable to image particular
solar regions ; instead they collected radiation from the
entire Sun. Using such an instrument, McKenzie &
Feldman (1992) derived the elemental abundance in a large
number of Ñares. According to them, Ñaring plasmas have
FIP biases that vary between 1 and 4. In general, the McKenzie & Feldman (1992) results were conÐrmed by most
subsequent studies, although in few cases somewhat di†erent conclusions were made (for more details see the above
referenced review papers). SUMER is the Ðrst instrument
providing sufficient number of lines that could be used to
determine abundances of well-deÐned high-temperature
solar plasmas.
A very important issue regarding elemental abundances
is the dependence of the FIP bias on the FIP, i.e., do all
low-FIP elements get enriched in equal proportions or are
they being enriched in proportion to the magnitude of their
FIPs ? The high-temperature SUMER spectra contains forbidden lines from elements spanning a wide FIP range, in
particular, lines emitted by the following ions : Ar10`,
520
FELDMAN ET AL.
Vol. 544
TABLE 11
N-LIKE LINES WITHIN 2s22p3 GROUND CONFIGURATION IN Ar XII, Ca XIV, AND Fe XX SUITABLE FOR
SPECTROMETER CALIBRATIONS RESPONSIVITY OVER WIDE VACUUM WAVELENGTH RANGES
Ar XII
UPPER
LEVEL
LOWER
LEVEL
2P . . . . . .
3@2
2P . . . . . .
3@2
2P . . . . . .
3@2
2P . . . . . .
3@2
2P . . . . . .
1@2
2P . . . . . .
1@2
Ca XIV
Fe XX
j (AŽ )
Aa (s~1)
j (AŽ )
Aa (s~1)
j (AŽ )
Aa (s~1)
2P
1@2
2D
5@2
2D
3@2
4S
3@2
20585b
1789.29b
1688.76b
649.093
0.924
240
476
1170
9138b
1432.228
1291.607
545.225
10.1
725
1620
3230
1586.309
679.269
541.35
309.32
1600
12800
44900
27300
2D
3@2
2S
3@2
1839.69b
670.302
241
549
1504.275
579.853
666
1730
821.706
384.22
6170
31500
0.377
2.17
13153b
880.401
4.19
13.5
2665.1
567.825
417
1240
2D . . . . . .
2D
30057b
5@2
3@2
2D . . . . . .
4S
1018.790
5@2
3@2
a Spontaneous decay rate.
b Line not yet observed.
Ar11`, Ar12` (FIP \ 15.8 eV) ; Fe11` (FIP \ 7.9 eV) ;
Ni12`, Ni13`, Ni14` (FIP \ 7.6 eV) ; Ti14` (FIP \ 6.8
eV) ; Cr15`, Cr16` (FIP \ 6.7 eV) ; Ca11`, Ca12`, Ca13`,
Ca14` (FIP \ 6.1 eV) ; and K12` (FIP \ 4.3 eV). As was
described in ° 6, lines from the Ca ions could be used to
derive emission measure distribution over the 1 ] 106 \
T \ 5 ] 106 K. Similarly, the electron density and teme
perature
in the emitting plasma could also be derived using
the available SUMER line ratios (see °° 5 and 7). With the
emission measure, electron density, and temperature
known, the elemental abundance relative to Ca of the high
FIP Ar, the low FIP Fe, Ni, and the very low FIP K could
be derived. Similarly, as a check on the determination, elemental abundances of Cr and Ti that have FIP similar to
Ca could also be derived. In doing so the two questions of
““ What are the elemental abundances in Ñares ? ÏÏ and ““ Do
FIP bias levels depend on the size of the FIP ? ÏÏ could be
answered.
Previous attempts to determine the dependence of the
FIP bias levels on the size of the FIP were attempted.
Although they appeared to show such a trend, the results
were inconclusive (Feldman 1992).
9. SETS OF SPECTRAL LINES ORIGINATING FROM
COMMON UPPER LEVELS WITH VASTLY DIFFERENT
WAVELENGTHS
The responsivity calibration of a space spectrometer that
operates over a wide wavelengths range is quite difficult.
Even in cases when the responsivity as a function of wavelength could be established on the ground, once in space the
calibration needs to be checked or reestablished from time
to time. A common technique of such a calibration utilizes
sets of lines with known branching ratios that originate
from the same upper level and have largely di†erent wavelengths. For such lines to be useful, their branching ratios
should not be too di†erent, resulting in intensity ratios that
are not too di†erent. Unfortunately, the task of Ðnding sets
of lines fulÐlling the above requirement is often difficult
because in most cases allowed lines of comparable branching ratios appear in close wavelength proximities.
An interesting case are the nine forbidden lines that result
from transitions between the 4S , 2D , 2D , 2P , and
3@2 conÐguration
3@2
5@2 in1@2N-like
2P
levels of the 2s22p3 ground
3@2
ions. As a result of the unique distribution of ““ j ÏÏ values
within levels of the ground conÐguration, the highest level
2P decays into all four lower levels, resulting in lines of
3@2di†erent wavelengths. The second highest level 2P
very
1@2
decays into the 2D , and 4S , and the third highest level
3@2
3@2
2D
level decays into either 2D , and 4S . Table 11
5@2 the wavelengths and spontaneous
3@2
gives
decay3@2rates of forbidden lines that arise from levels within the ground conÐguration of Ar11`, Ca13` and Fe19` that could be used to
calibrate the efficiency of spectrographs spanning a very
wide wavelength range. Mason & Bhatia (1983) and
Kaufman & Sugar (1986) published spontaneous decay
rates for the Fe XX forbidden transitions. The agreement
between the two sets of calculations was better than 15%.
Recently, Bhatia (1999, private communication) recalculated the Fe XX spontaneous decay rates using the latest
experimental energy level values. His newly calculated
values that almost always fell in between values from the
previous two calculations are shown in the last column of
TABLE 12
C-LIKE LINES WITHIN THE 2s22p3 GROUND CONFIGURATION IN Ar XIII, Ca XV, AND Fe XXI
SUITABLE FOR SPECTROMETER RESPONSIVITY CALIBRATIONS
Ar XIII
UPPER
LEVEL
LOWER
LEVEL
j (AŽ )
1D . . . . . .
3P
1330.532
2
1
1D . . . . . .
3P
1583.36b
2
2
a Spontaneous decay rate.
b Line not yet observed.
Ca XV
Fe XXI
Aa (s~1)
j (AŽ )
Aa (s~1)
j (AŽ )
Aa (s~1)
150
266
1098.484
1375.959
551
810
585.766
786.031
15900
15100
No. 1, 2000
SPECTRAL LINES AND THEIR POTENTIAL USE IN PLASMA DIAGNOSTICS
Table 11. The spontaneous decay rates of Ar XII and Ca XIV,
taken from Kaufman & Sugar (1986), are also given in the
table.
The 2s22p2 ground conÐguration in C-like ions and the
2s22p4 ground conÐguration in O-like ions each contains
line pairs that arise from transitions between 1D ] 3P
2
2
and 1D ] 3P that in principle could be used for efficiency
2
1
calibration. Unfortunately, the 1D È3P transitions in
O-like Ar XI, Ca XIII, and Fe XIX are2very1weak and probably would not be useful in most instances. Table 12 provides the wavelengths and spontaneous decay rates of the
set of lines in the C-like Ar, Ca, and Fe ions that are useful
for efficiency calibration.
The wavelengths displayed in Table 11 span a very wide
range. The Fe XX lines span the 300È2660 AŽ range, Ca XIV
span the 550È13000 AŽ range, and the Ar XII lines span 650È
30000 AŽ . All together the table displays 24 lines that are
divided among six pairs and three sets of four lines each that
521
together span the 300È30000 AŽ wavelength range. Notice
that the ratios of the spontaneous decay rates within each
pair and within three lines from the sets of four are within 1
order of magnitude of each other. As a result, the number of
photons that would be emitted in lines with vastly di†erent
wavelengths that belong to the same set would not be too
di†erent.
The work of U. F. was supported by the ONR/NRL
Research Option, Solar Magnetism and EarthÏs Environment, and by NASA SR&T grant. W. C., E. L., and K. W.
were supported by the MPG. SUMER is Ðnancially supported by DLR, CNES, NASA, and the ESA PRODEX
program (Swiss contribution). SOHO is a project of international cooperation between ESA and NASA. We also
acknowledge the comments we received from B.N. Dwivedi,
who critically reviewed this paper.
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