THE ASTRONOMICAL JOURNAL, 117 : 1505È1548, 1999 March
( 1999. The American Astronomical Society. All rights reserved. Printed in U.S.A.
A GODDARD HIGH RESOLUTION SPECTROGRAPH ATLAS OF ECHELLE OBSERVATIONS OF THE
HgMn STAR s LUPI
J. C. BRANDT,1,2 S. R. HEAP,1,3 E. A. BEAVER,1,4 A. BOGGESS,1,5 K. G. CARPENTER,1,3 D. C. EBBETS,1,6 J. B. HUTCHINGS,1,7
M. JURA,1,8 D. S. LECKRONE,1,3 J. L. LINSKY,1,9 S. P. MARAN,1,10 B. D. SAVAGE,1,11 A. M. SMITH,1,3 L. M. TRAFTON,1,12
F. M. WALTER,1,13 R. J. WEYMANN,1,14 C. R. PROFFITT,15,16 G. M. WAHLGREN,15,16 S. G. JOHANSSON,17 H. NILSSON,17
T. BRAGE,17 M. SNOW,15,2 AND T. B. AKE15,18,16
Received 1998 February 20 ; accepted 1998 December 10
ABSTRACT
Observations of the ultraÈsharp-lined, chemically peculiar star s Lupi taken by the Goddard High
Resolution Spectrograph in echelle mode are presented. Thirty-six intervals of the spectral region
between 1249 and 2688 AŽ are covered with resolving powers in the range 75,000È93,000. Line identiÐcations are provided, and the observed spectra are compared with synthetic spectra calculated using the
SYNTHE program and associated line lists with changes to the line lists. The signiÐcance of these
spectra for the s Lupi PathÐnder Project and the closely related atomic physics e†ort is discussed in a
companion paper.
Key words : atlases È stars : chemically peculiar È stars : individual (s Lupi)
1.
INTRODUCTION
fully characterized observationally and which are among
the most puzzling aspects of CP star spectra. An excellent
overview of the various classes of CP stars may be found in
Wol† (1983) and in Smith (1996).
Most of the modern interpretation of CP star spectra is
based on the mechanism of radiatively driven di†usion
within the stellar atmosphere or envelope by which individual elements and their ions are segregated, stratiÐed, blown
away from the star, or left unsupported to sink below the
visible surface layers (Michaud 1970). Radiatively driven
di†usion can lead to signiÐcant anomalies, however, only
within a highly stable hydrodynamic environment. As
observational testimony to this stability, some of the CP
stars, including the subject of this paper, exhibit the sharpest spectral lines seen on the upper main sequence.
In the modern era of quantitative abundance analyses of
stellar spectra with high-speed computers and sophisticated
model atmospheres, the critical testing and assessment of
theoretical scenarios for the production of a broad range of
abundance and isotopic anomalies have been severely
limited by the sparsely populated optical-wavelength
spectra of late B- and early A-type dwarfs. At wavelengths
accessible from ground-based telescopes, spectroscopic
observations o†er only sketchy coverage of the periodic
table, usually providing information about only one ionization state of an element and sometimes giving very misleading abundance results. The obvious solution to this
quandary resides in high-resolution observations in the
satellite ultraviolet, where the spectra of CP stars are
extraordinarily rich in low-excitation lines of the Ðrst two or
three ionization states of many elements.
With the launch of the Goddard High Resolution
Spectrograph (GHRS) on the Hubble Space T elescope
(HST ) in 1990 (Brandt et al. 1994), a tool of remarkable
power became available for the Ðrst detailed, highresolution exploration of the ultraviolet spectrum of a CP
star. The primary target selected for this endeavor is the
ultraÈsharp-lined, nonmagnetic star, s Lupi (B9.5p
HgMn ] A2 Vm). For the past 8 years an international
team of astrophysicists and atomic physicists, led by D. S.
Leckrone, has carried out a systematic exploration of the
exotic ultraviolet spectroscopic ““ landscape ÏÏ of this arche-
Magnetic and nonmagnetic chemically peculiar (CP)
stars make up approximately 10%È20% of the total Galactic population of main-sequence stars with e†ective temperatures between 9500 and 16,000 K. Their photospheric
abundance anomalies vary widely in magnitude from
element to element across the entire periodic table and from
star to star. In some cases these anomaliesÈenhancements
and depletions relative to normal stellar abundancesÈare
of enormous magnitude. They include anomalous isotopic
compositions of a number of elements, which are not yet
ÈÈÈÈÈÈÈÈÈÈÈÈÈÈÈ
1 GHRS Investigation DeÐnition team.
2 Laboratory for Atmospheric and Space Physics, Campus Box 392,
University
of
Colorado,
Boulder,
CO
80309-0392 ;
brandt=lyrae.colorado.edu.
3 Laboratory for Astronomy and Solar Physics, Goddard Space Flight
Center, Code 681, Greenbelt, MD 20771.
4 Center for Astrophysics and Space Sciences, University of California,
San Diego, C-0111, La Jolla, CA 92093-0111.
5 2420 Balsam Drive, Boulder, CO 80304.
6 Ball Aerospace and Technologies Corporation, P.O. Box 1062, AR1,
Boulder, CO 80306.
7 Dominion Astrophysical Observatory, 5071 West Saanich Road, Victoria, V8X 4M6 BC, Canada.
8 Department of Physics and Astronomy, University of California, Los
Angeles, Los Angeles, CA 90095-1562.
9 JILA, University of Colorado and National Institute of Standards
and Technology, Boulder, CO 80309-0440.
10 Space Sciences Directorate, Code 600, Goddard Space Flight Center,
Greenbelt, MD 20771.
11 Department of Astronomy, University of Wisconsin, 475 North
Charter Street, Madison, WI 53706.
12 MacDonald Observatory and Astronomy Department, University of
Texas, Austin, TX 78712.
13 Astronomy Program, Department of Earth and Space Sciences, State
University of New York at Stony Brook, Stony Brook, NY 11794-2100.
14 Observatories of the Carnegie Institution of Washington, 813 Santa
Barbara Street, Pasadena, CA 91101.
15 GHRS Science team.
16 Science Programs, Computer Sciences Corporation, Goddard Space
Flight Center, Code 680, Greenbelt, MD 20771.
17 Department of Physics, University of Lund, Box 118, S-22100 Lund,
Sweden.
18 Currently at Department of Physics and Astronomy, Johns Hopkins
University, Baltimore, MD 21218.
1505
1506
BRANDT ET AL.
Vol. 117
TABLE 1
HIGH S/N GHRS SSA ECHELLE OBSERVATIONS OF s LUPI
j
min
(AŽ )
j
max
(AŽ )
Root Name
Figure
Echelle Order
t
exp
(s)
Maximum S/N
Resolution
(AŽ )
j [j
A
B
(AŽ )
1249.03 . . . . . .
1300.79 . . . . . .
1317.75 . . . . . .
1331.48 . . . . . .
1356.91 . . . . . .
1373.50 . . . . . .
1409.43 . . . . . .
1418.82 . . . . . .
1433.33 . . . . . .
1507.09 . . . . . .
1535.33 . . . . . .
1552.64 . . . . . .
1644.73 . . . . . .
1736.50 . . . . . .
1844.31 . . . . . .
1862.87 . . . . . .
1863.38 . . . . . .
1904.60 . . . . . .
1936.74 . . . . . .
1936.82 . . . . . .
1997.73 . . . . . .
2012.17 . . . . . .
2023.00 . . . . . .
2059.71 . . . . . .
2137.18 . . . . . .
2147.73 . . . . . .
2201.25 . . . . . .
2262.83 . . . . . .
2271.39 . . . . . .
2323.93 . . . . . .
2335.20 . . . . . .
2346.95 . . . . . .
2376.15 . . . . . .
2407.62 . . . . . .
2434.82 . . . . . .
2528.21 . . . . . .
2602.06 . . . . . .
2673.98 . . . . . .
2673.78 . . . . . .
2674.51 . . . . . .
1255.67
1307.82
1324.50
1338.69
1364.41
1380.72
1416.82
1426.83
1441.10
1515.33
1543.11
1561.06
1653.64
1746.27
1853.44
1873.10
1873.62
1915.57
1947.06
1947.23
2008.66
2022.72
2033.45
2071.23
2149.17
2159.53
2212.05
2274.57
2282.97
2336.75
2347.84
2359.37
2387.90
2421.18
2448.04
2542.33
2614.81
2688.32
2688.11
2688.85
Z2MB0107
Z2MB0108
Z2MB010B
Z2MB010C
Z2MB010D
Z2MB010G
Z2MB010H
Z3650308
Z2MB010I
Z365030B
Z2MB010L
Z3650307
Z2MB010M
Z0IX010R
Z0IX010J
Z13J010M
Z13J510M
Z28H010M
Z0G7010J
Z0IX010B
Z16C0107
Z16C0108
Z16C010B
Z16C010C
Z16C010F
Z16C010G
Z16C010J
Z16C010K
Z16C010N
Z16C010O
Z16C010R
Z0IX010O
Z16C010S
Z16C010V
Z16C010W
Z0IX010G
Z16C010X
Z28H010L
Z13J010L
Z13J510L
1
2
3
4
5
6
7
8
9
10
11
12
13
14
15
16
...
17
18
...
19
20
21
22
23
24
25
26
27
28
29
30
31
32
33
34
35
36
...
...
A45
A43
A43
A42
A41
A41
A40
A39
A39
A37
A37
A36
A34
B32
B31
B30
B30
B29
B29
B29
B28
B28
B28
B27
B26
B26
B26
B25
B25
B24
B24
B24
B24
B23
B23
B22
B22
B21
B21
B21
2585.09
2154.24
2908.22
1507.97
1723.39
1507.97
1723.39
2369.66
1723.39
3015.94
2154.24
3191.56
3115.73
5170.18
2585.09
2585.09
2585.09
2585.09
2423.52
2154.24
1077.12
1077.12
1077.12
969.41
753.98
646.27
1077.12
538.56
753.98
646.27
646.27
861.70
969.41
538.56
538.56
861.70
753.98
646.27
861.70
861.70
56
73
87
76
75
76
75
81
75
94
59
88
68
47
31
85
85
100
95
52
71
63
67
74
72
60
67
66
75
67
79
53
69
58
78
58
63
141
103
109
0.014
0.015
0.015
0.016
0.016
0.016
0.016
0.017
0.017
0.017
0.017
0.018
0.019
0.020
0.021
0.022
0.022
0.022
0.023
0.023
0.024
0.023
0.023
0.027
0.026
0.025
0.024
0.025
0.025
0.028
0.027
0.028
0.026
0.030
0.029
0.030
0.028
0.031
0.031
0.031
0.253
0.249
0.246
0.233
0.225
0.215
0.215
[0.205
0.205
[0.248
0.213
[0.204
0.211
[0.451
[0.442
[0.743
0.417
[0.483
[0.799
[0.425
[0.357
[0.354
[0.339
[0.340
[0.333
[0.330
[0.318
[0.321
[0.300
[0.302
[0.278
[0.590
[0.279
[0.255
[0.239
[0.585
[0.246
[0.682
[1.077
0.593
typal CP star. The data set consists of 36 wavelength intervals, averaging 10 AŽ in width, observed in the GHRS echelle
mode with resolving powers ranging from 75,000 to 93,000.
These observations were obtained over the course of 32
HST orbits.
This project has been named the ““ s Lupi PathÐnder
Project ÏÏ because it is entirely analogous to the exploration
and mapping of a previously unexplored land. As described
in a companion paper (Leckrone et al. 1999, hereafter Paper
I) and in an earlier review (Leckrone et al. 1998), the
analysis of the ultraviolet spectrum, supplemented by a new
study of s LupiÏs optical-wavelength spectrum (Wahlgren,
Adelman, & Robinson 1994), has resulted to date in a quadrupling of the number of elements for which accurate abundances or upper limits to abundances in the photosphere of
this star are known. The s Lupi PathÐnder Project has also
identiÐed previously unknown isotope anomalies and has
produced several new lines of qualitative and quantitative
evidence bearing on the mechanism of radiatively driven
di†usion.
A no less important result has been the establishment of
s Lupi as an internationally accepted astrophysical
““ standard light source ÏÏ for atomic spectroscopy. Its ultraviolet spectrum displays a broad panorama of transitions of
both rare and abundant elements, which are difficult to
observe directly in the laboratory, and the GHRS observations exceed in resolution and wavelength accuracy all
but the very best (Fourier Transform Spectrometer) laboratory measurements. The direct participation of atomic
physicists in the exploration and analysis of s LupiÏs ultraviolet spectrum has been essential, given the inadequacy of
the existing body of basic atomic data in the literature. In
turn, the astrophysical data have provided new insights
about atomic structure (e.g., Johansson et al. 1995 ; Leckrone et al. 1996a).
The observations of s Lupi were part of the GHRS Guaranteed Time Observer (GTO) program. The Investigation
DeÐnition Team (IDT) agreed to pursue several ““ team
projects ÏÏ that would utilize the unique capabilities of the
GHRS to produce atlases of astrophysically important
No. 3, 1999
ATLAS OF SPECTRUM s LUPI
objects. The Ðrst Ðve of these dealt with 3C 273 (Brandt et
al. 1993), a Orionis (Brandt et al. 1995), f Ophiuchi (Brandt
et al. 1996), 3C 273 again (Brandt et al. 1997), and 10
Lacertae (Brandt et al. 1998).
In the present paper we provide a complete atlas of the
GHRS FP-SPLIT echelle observations of s Lupi, together
with the best synthetic spectrum the s Lupi PathÐnder
Project team is currently able to provide for the case of a
homogeneous photosphere in LTE, given the present state
of knowledge of elemental abundances and atomic parameters. The study of this extraordinary spectrum is far from
complete. We anticipate that the atlas will be a resource for
both astrophysics and atomic physics and will serve as a
road map for future ultraviolet spectroscopy of a variety of
early-type stars.
2.
DATA REDUCTION
All data were collected using the FP-SPLIT \ 4 option,
which divides the observation into four parts taken at
slightly di†erent positions of the grating carrousel. This
shifts the spectrum along the photocathode and allows
blemishes and diode irregularities to be distinguished from
intrinsic spectral features. Data points at a given pixel that
systematically di†ered from the median Ñux at a Ðxed wavelength by more than a few standard deviations in one-half
or more of the subexposures were identiÐed as such irregularities and excluded from the Ðnal co-addition. This procedure eliminated obvious blemishes while only rarely
excluded apparently good data points. The data were otherwise reduced using the standard CALHRS software. A
detailed description of this procedure can be found in Wahlgren et al. (1995).
Table 1 summarizes all GHRS small science aperture
(SSA) echelle observations of s Lupi that used the FPSPLIT \ 4 option for the reduction of Ðxed pattern noise.
For a few wavelength intervals, duplicate observations
exist. These are listed in Table 1, but Ðgures displaying these
observations are not presented here. A few large science
aperture echelle observations with FP-SPLIT \ 4 also
exist, but in each case these merely duplicate the wavelength
coverage of an existing SSA observation.
The minimum and maximum wavelengths listed for each
observation in Table 1 are for the rest velocity of s Lupi A.
The maximum single-pixel signal-to- noise ratio (S/N) for
each spectrum, as calculated using Poisson statistics, is also
given. Due to the incomplete removal of low-level Ðxed
pattern noise and variations in sensitivity with wavelength,
this only provides an upper limit to the true continuum
S/N.
3.
SPECTRUM MODELING
The model atmospheres of the primary and secondary
stars that we use for spectral synthesis calculations are the
same ATLAS8 models described by Wahlgren et al. (1994).
Synthetic spectra were calculated using the program
SYNTHE and associated line lists from Kurucz (1991,
1993). Changes to the line lists of Kurucz are described
below and in Paper I.
The wavelength shift between the primary and secondary
stars for the start time of each observation was calculated
using the orbital parameters determined by Dworetsky
(1972). Each synthetic spectrum was rotationally broadened, and they were then co-added using the calculated
wavelength shift and assuming that the Ñux ratio of the two
1507
stars is given by H (l)R2/H (l ] *l)R2, where H(l) is the
A
A B
B
surface brightness per unit surface area at each frequency
calculated by SYNTHE and the ratio of the radii R /R \
A B
1.67 as determined by Wahlgren et al. (1994). Each combined spectrum was then convolved with a Gaussian of
width equal to the expected instrumental broadening for
that observation. The e†ects of the secondary spectrum are
most clearly seen in Figure 32 below, where several features
in the observed spectrum are clearly due to lines in the
secondary (e.g., Co II j2408.414 shifted to j2408.16 and Fe II
j2409.379 shifted to j2409.13). The observed shifts for these
secondary lines are within 2 km s~1 (0.016 AŽ ) of their predicted values.
4.
DESCRIPTION OF THE FIGURES
In each Ðgure, the solid line shows the observed GHRS
spectrum after aligning and merging individual subexposures and removing data a†ected by blemishes. The
dotted line shows the combined synthetic spectra for both
the primary and the secondary stars, including both rotational and instrumental broadening. The dashed line shows
the synthetic spectrum for the secondary alone, also with
rotational and instrumental broadening included. Note that
the line for the synthetic spectrum of the secondary cannot
be seen in Figure 1 below, because of the low Ñux level for
the secondary calculated at that wavelength. The various
symbols annotating the line identiÐcations are discussed in
the following text and are summarized in Table 2.
The labeled wavelength scale is aligned with the laboratory wavelengths used in the synthetic spectrum of the
primary star. Vacuum wavelengths are given below 2010 AŽ ,
and air wavelengths, above this. The observed GHRS data
have been shifted in wavelength to align with the combined
synthetic spectra. We have only adjusted the zero-point
o†set of the wavelength calibration, although in some cases
a small change in the dispersion might have improved the
Ðt. A description of the default wavelength calibration procedure for the GHRS can be found in Robinson et al. (1992).
The horizontal arrow below the second and fourth panels
of each Ðgure shows the amount and direction that the
secondaryÏs wavelength is shifted with respect to the labeled
wavelength scale. In Figure 3 below, for instance, the dip in
the secondary spectrum at 1319.75 AŽ in the labeled wavelength scale corresponds to 1319.5 AŽ at the rest wavelength
of the secondary. This dip in the secondary spectrum, thus,
is caused by the same blend of Fe II lines seen in the primary
near 1319.5 AŽ .
Many wavelengths have been updated using Fourier
Transform Spectrometer (FTS) data from the University of
Lund or from Imperial College, London. These wavelengths
should be especially accurate, and we have marked these
wavelengths with an ““ f ÏÏ but have not usually provided
further reference information. Note that the task of updating the observed transitions with FTS data is far from comTABLE 2
ANNOTATIONS OF LINE IDENTIFICATIONS
Symbol
Meaning
* .......
¤ .......
” .......
f ......
A ......
Discussed in the notes for that Ðgure or in Paper I
Kurucz ““ guess ÏÏ for gf value
Arbitrary modiÐcation to gf value
j from Fourier transform spectroscopy
““ Superanomalous ÏÏ transition
1508
BRANDT ET AL.
plete. Many additional wavelengths could be improved by
comparison with existing laboratory data, but care must be
taken to ensure correct identiÐcation of laboratory lines
with stellar features, as several laboratory transitions may
be close to an observed spectral line. A purely automated
approach would lead to misidentiÐcation of some transitions.
Features in the synthetic spectrum corresponding to
absorption lines in the primary are labeled in a plain text
font. To aid in line identiÐcation, the fractional part of the
wavelength is given next to the ion name. Ion names are
marked with an asterisk if there is additional discussion
or information regarding this line either in this paper or in
Paper I. If an alternative wavelength or transition probability has been used in place of data from the Kurucz line lists,
this is detailed in Tables 1 and 2 of Paper I, although wavelength updates from Lund or Imperial College FTS measurements are not included there. Other information is
given in the notes on individual observations (° 5).
The line identiÐcations were based on analytic estimates
of line strengths made using the SYNSPEC code (version
38) of Hubeny & Lanz (1995). The threshold for which lines
to label was adjusted for each spectrum. Lines that clearly
make only minor contributions to a blend were not labeled
even if above this threshold, while weaker lines that correspond to an obvious feature in the synthetic spectrum were
labeled. Because the analytic estimates can be misleading,
especially when comparing lines formed at very di†erent
depths, we have not listed them on the Ðgures.19
It should be emphasized that the synthetic spectra are
intended to be only a guide to the data presented in this
atlas and do not constitute a deÐnitive interpretation. For
example, di†erent ions of the same element may yield very
di†erent abundances. In some cases, this is likely to be a
non-LTE e†ect that cannot be modeled with the LTE spectrum synthesis presented here. For other ions it may reÑect
inaccurate transition probabilities, inadequacies in the
model atmosphere, or vertical stratiÐcation of chemical
abundances.
The extensive abundance analyses carried out to date as
part of the s Lupi PathÐnder Project are summarized in
Paper I. Where inconsistent abundances have been determined from di†erent ionization stages of the same element,
we adopt for this atlas the abundance determined from lines
of the dominant ionization stage in the atmosphere of s
Lupi. This may result in very poor Ðts in the atlas to lines of
the other ions.
This is especially apparent for the lines of many neutral
species that are systematically weaker than predicted when
using abundances determined from singly charged ions.
Such anomalies clearly exist for C I and Hg I lines. Similar
anomalies are also evident for Si I and P I, although for
these ions the largest discrepancies appear to be at least
partly due to inaccurate transition probabilities. Any deÐnitive analysis of ionization ratios in these elements will need
to be preceded by a critical evaluation of the available transition probability data. Some examples of such discrepancies will be mentioned in the notes on individual
observations. A similar e†ect is also seen for a number of
ÈÈÈÈÈÈÈÈÈÈÈÈÈÈÈ
19 This paperÏs referee suggested that readers may Ðnd the VALD database (Piskunov et al. 1995) useful for making their own estimates of relative
line strengths.
Vol. 117
doubly charged ionsÈe.g., Zr II j1938.500 versus Zr III
jj1937.215, 1940.236, and 1941.053 in Figure 18 below, or
Y II j2413.913 versus Y III j2414.643 in Figure 32 below.
Where a known laboratory line coincides with an
unmodeled feature in the observed data, we have marked
the line using an italic font. However, these should be considered as suggested identiÐcations only. Further analysis of
the atomic data and stellar spectrum is needed to conÐrm or
refute these suggestions.
In many cases, the synthetic spectrum will show a strong
line that is weak or absent in the observed data or will
dramatically underestimate the strength of an observed line.
For many elements this reÑects the varying quality of the
input line lists. Lines marked with a dagger denote oscillator strengths that Kurucz references as ““ guessed.ÏÏ These
are really values inserted as ““ placeholders ÏÏ where multiplet
tables indicated a line but no transition probability was
available and were never intended to approximate the true
transition probabilities. For some other lines in the Kurucz
lists, the oscillator strengths given there are estimates based
on the relative intensities of laboratory emission line
spectra ; these transition probabilities often su†er from large
systematic errors.
We have, in general, avoided making arbitrary changes to
transition probabilities simply to improve the agreement
between the observed and synthetic spectra. However, in
some cases anomalously strong calculated features
obscured other relevant details in the synthetic spectrum.
We then made arbitrary reductions in certain transition
probabilities to prevent this. Lines altered for this reason
are marked with a double dagger.
Large discrepancies also occur for some iron group lines
calculated by Kurucz. As was discussed by Leckrone et al.
(1993a), such ““ superanomalous ÏÏ iron group lines can occur
for a number of reasons. Energy levels from KuruczÏs
Cowan code (Cowan 1981) calculations are sometimes
matched with the wrong experimentally measured energy
levels. Published term analyses sometimes contain errors.
Even when levels are identiÐed correctly, accidental coincidence between energy levels can lead to strong conÐguration interactions and level mixing. These e†ects can be
sufficiently sensitive to the exact level placement that no ab
initio theoretical code can avoid large errors in calculated
oscillator strengths. Some particularly extreme examples of
such level mixing in Fe II were discussed by Johansson et al.
(1995), and several of these lines are discussed in the notes
on individual observations. We have noted several examples of ““ superanomalous ÏÏ lines of iron group elements that
are predicted to be much stronger than indicated by the
GHRS data by marking them with an ““ A ÏÏ in the Ðgures.
Having pointed out some of the inaccuracies to be found
in a small minority of the listings in the Kurucz database,
we hasten to note that this database is the indispensible
foundation on which all modern spectrum synthesis calculations rest. No other set of atomic parameters for individual spectral transitions is so large, nor, on average, so
accurate. The analysis of a complex stellar spectrum such as
that of s Lupi could not proceed without the Kurucz database as the starting point.
A few observations show evidence for interstellar lines.
This is most apparent in Figures 29 and 31 below, in which
interstellar components of Fe II jj2343.495 and 2382.038
are visible. For these and other cases in which interstellar
absorption is suspected, we have marked the location of
No. 3, 1999
ATLAS OF SPECTRUM s LUPI
components shifted by [6 and [12 km s~1 from the
center-of-mass velocity of the binary. Dworetsky (1972) suggested that there were long-term variations in the system
velocity of a few km s~1, indicating the possible presence of
a third body, and so it cannot be assumed that observations
taken at di†erent epochs will always show the interstellar
features at the same displacement from the center-of-mass
velocity of the two conÐrmed components. The data shown
here are not adequate to address this question, as all of the
observations in this data set showing clearly resolved interstellar lines (Figs. 29 and 31 and possibly Figs. 21 and 22, all
below) were taken within an 8 hour period on 1993 February 1 and 2. However, a low S/N calibration observation
taken without using the FP-SPLIT option (observation
Z3D5041IT from 1996 August 17) appears to show interstellar Al II components shifted by [3 and [9 km s~1 from
the center-of-mass velocity. This 3 km s~1 change from the
1993 observations supports DworetskyÏs suggestion of
small long-term velocity shifts.
5.
NOTES ON INDIVIDUAL OBSERVATIONS20
Observation Z2MB0107, 1249.03È1255.67 AŽ , Figure 1.È
Especially poor Ðts are found for the S I lines at 1253.297
and 1253.325 AŽ .
For the Si II lines, we substituted the semiclassical Kurucz
& Peytremann (1975) damping coefficients for the values
given by Kurucz. This especially improved the Ðt for the
Si II j1251.164 line. The continuum edge for the ground
state of Si I is near 1520 AŽ , so it is not surprising that the
proÐle of the autoionizing line at 1255.272 AŽ is poorly
reproduced by the synthetic spectrum. The Ðt could be
improved by drastically increasing the damping width of
the line.
This region is also predicted to contain numerous C I
lines for which laboratory data are not available (cf.
KuruczÏs line lists). These lines were not included in our
spectral synthesis as their true wavelengths are unknown,
but they are expected to produce discrete absorption features in the stellar spectrum. These lines have the 2p2 1D
level at an excitation energy of 10192 cm~1 as their lower
level. The Ðrst continuum limit for this level is near 1240 AŽ .
Observation Z2MB0108, 1300.79È1307.82 AŽ , Figure
2.ÈThe spectrum in this region is dominated by a number
of very strong Si II lines. For three Si II lines from excited
levels (those marked with a double dagger) we have arbitrarily reduced the transition probabilities taken from
Kurucz by 0.5 dex to reduce the mismatch between the
synthetic and the observed spectra. The strongest of these
lines (1305.592 AŽ ), is an autoionization line for which Artru
et al. (1981) give a transition probability 0.18 dex smaller
than that of Kurucz. Using this value would still produce
too much absorption ; however, the broadening parameters
are not well determined.
The wavelength for the Tl II line at 1307.50 AŽ (Ellis &
Sawyer 1936) is of limited accuracy, as is the Cowan code
calculation by Brage used for the transition probability
(Paper I, Table 2). Either of the adjacent unidentiÐed features in the observed spectrum might be the thallium line,
or if there is substantial hyperÐne structure, both features
might be due to this transition.
ÈÈÈÈÈÈÈÈÈÈÈÈÈÈÈ
20 The reduced data and the synthetic spectra may also be found on the
World Wide Web in digital form at http ://archive.stsci.edu/.
1509
The observed S I lines (1302.336, 1302.862, 1303.430, and
1305.884 AŽ ) are uniformly weaker than predicted. While
this may suggest an ionization anomaly, similar discrepancies appear for some S II lines (see the notes for Fig. 11).
We have marked the location of interstellar components
shifted by [6 and [12 km s~1 from the center-of-mass
velocity for the ground-state lines of O I 1302.168 AŽ and Si II
1304.370 AŽ , corresponding to the components seen in the
Fe II line at 2343.495 AŽ in Figure 29. The orbital phase at
which this observation was taken shifts the stellar line proÐles on top of the possible interstellar components, but both
lines clearly show excess absorption and sharp-edged proÐles on the long wavelength sides.
Observation Z2MB010B, 1317.75È1324.50 AŽ , Figure
3.ÈA preliminary version of this Ðgure was published in
Proffitt et al. (1998). Noteworthy in this interval are the two
hyperÐne components of the Tl II resonance line near
1321.64 AŽ . The synthetic spectrum assumes the same
isotope mixture (pure 205Tl) and abundance found by Leckrone et al. (1996b) from the ground-state intercombination
line near 1908 AŽ , (Fig. 17).
The transition probability for the Hg II line at 1321.719 AŽ
(log gf \ [0.005) is from the Cowan code calculation of
Brage et al. (1999), rather than from their MultiConÐguration DiracÈFock calculations. See Proffitt et al.
(1999) for further discussion of this line. All mercury lines in
this atlas are calculated assuming an isotope blend of 1%
202Hg and 99% 204Hg (White et al. 1976 ; Proffitt et al.
1998), and the listed wavelengths are those of the dominant
isotope, 204Hg.
Observation Z2MB010C, 1331.48È1338.69 AŽ , Figure
4.ÈThe overall carbon abundance adopted for these
models was adjusted to yield an approximate Ðt to the
strong C II lines (1334.532, 1335.663, and 1335.708 AŽ ) in this
wavelength interval. The line proÐles calculated here appear
slightly too strong in the line cores and slightly too weak in
the line wings. A possible interstellar feature in the red wing
of the 1334.532 AŽ line is also noted.
The transition probability for the Hg II line at 1331.747 AŽ
(log gf \ [0.020) is from the Cowan code calculations of
Brage et al. (1999).
Observation Z2MB010D, 1356.91È1364.41 AŽ , Figure 5.È
There is a clear carbon ionization anomaly in s Lupi, with
observed ultraviolet C I lines systematically weaker than
predicted by LTE synthetic spectra.
Limits on the B II abundance using the 1362.461 AŽ line
have been studied by Zethson et al. (1998). The B II line is
extremely weak or perhaps absent, and it is unresolved in an
apparent blend with an Fe II line at 1362.451 AŽ . This Fe II
line is computed to be much too strong if the gf value of
Kurucz [log gf \ [2.693] is adopted, and we have substituted log gf \ [3.5 to avoid obscuring the possible contribution of the boron line to the synthetic spectrum.
The strongest of the Cu II resonance lines is at 1358.773
AŽ . It is Ðtted quite well with the abundance [log N(Cu)/
N(H) \ [8.35] originally derived from the 1944.597 AŽ line
(Fig. 18) by Leckrone et al. (1993b). Other well-Ðt Cu II lines
are at 2148.984 AŽ (Figs. 23 and 24), and 2276.259 AŽ (Fig. 27).
Observation Z2MB010G, 1373.50È1380.72 AŽ , Figure
6.ÈThe As II lines in a G160M spectrum of this region were
discussed by Wahlgren et al. (1994) ; at that time this echelle
spectrum was not yet available. We have included hyperÐne
structure only for the line at 1375.073 AŽ (Brage & Leckrone
1995), but for clarity we have only noted one of the hyper-
FIG. 1.ÈObservation Z2MB0107. See the text for a description of line and symbol styles.
1510
FIG. 2.ÈObservation Z2MB0108
1511
FIG. 3.ÈObservation Z2MB010B
1512
FIG. 4.ÈObservation Z2MB010C
1513
FIG. 5.ÈObservation Z2MB010D
1514
FIG. 6.ÈObservation Z2MB010G
1515
1516
BRANDT ET AL.
Ðne components used in the synthetic spectrum. For further
discussion of the analysis of the As II lines see Paper I, ° 4.
The poor Ðt to the P I line at 1377.073 AŽ suggests a
possible ionization anomaly, but this cannot be conÐrmed
without a thorough understanding of the adopted transition probability.
Observation Z2MB010H, 1409.43È1416.82 AŽ , Figure
7.ÈA preliminary value of the gallium abundance has been
used to Ðt the wings of the resonance line at 1414.401 AŽ . The
absorption in the secondary due to Fe II j1414.163 is shifted
to 1414.378 AŽ . As this line is obviously absent from the
observed primary spectrum, we have removed this line from
the synthesis of the secondary to avoid interfering with the
Ga II line proÐle. However, there is clearly signiÐcant
opacity missing from the synthetic spectra for both the
primary and secondary stars in this feature.
Observation Z3650308, 1418.82È1426.83 AŽ , Figure
8.ÈThe apparent absence of the hyperÐne components of
the Bi III line (1423.308, 1423.356, 1423.480, 1423.529 AŽ )
suggests either that there is an ionization anomaly relative
to Bi II or that we have overestimated the bismuth abundance from the Bi II lines in Figure 9 below (see Paper I, ° 5).
Also noteworthy in this spectrum are a number of Ti III
lines ; two of which (1420.034 and 1421.755 AŽ ) appear to be
relatively unblended.
Observation Z2MB010I, 1433.33È1441.10 AŽ , Figure
9.ÈThe three hyperÐne components of the Bi II resonance
line (1436.777, 1436.821, and 1436.856 AŽ ) appear to be
present in this spectrum but are blended with other
unidentiÐed lines.
We have used the badly blended Pb II resonance line at
1433.905 AŽ to determine an upper limit for the lead abundance. This abundance is used in calculating the proÐles of
all the Pb II and Pb III lines covered by this atlas (see Paper
I, ° 5).
Observation Z365030B, 1507.09È1515.33 AŽ , Figure
10.ÈThe C I lines at 1510.981 and 1511.907 AŽ provide a
good illustration of the carbon ionization anomaly. Note
the resolved line of Sb II at 1513.241 AŽ , which yields the
antimony abundance adopted in Paper I.
Observation Z2MB010L , 1535.33È1543.11 AŽ , Figure
11.ÈIn the Kurucz line lists, an error in transcribing transition probabilities from the original reference (Hibbert
1988) led to the inclusion of transition probabilities for
several P II lines that are too large by a few orders of magnitude. We have substituted transition probabilities from
Morton (1991). The good Ðt to most of these P II lines (e.g.,
1535.923 and 1536.416 AŽ ) supports the abundance found by
Wahlgren et al. (1994) from optical observations.
The S II lines in this region (1535.869, 1540.738, and
1541.493 AŽ all are substantially weaker in the observed data
than in the synthetic spectrum.
Observation Z3650307, 1552.64È1561.06 AŽ , Figure
12.ÈThe hyperÐne components of the resonance line of
Tl III near 1557.8 AŽ do not appear where predicted in this
interval (see Paper I, ° 5), but we suspect this results from an
error in the laboratory wavelengths of Joshi & Raassen
(1990).
We have marked the location of the two hyperÐne components of the 207Pb III ground-state intercombination line
(1553.006 and 1553.051 AŽ ). The components from the even
isotopes of lead lie between these. The observed line near
1553.004 AŽ can be reasonably well Ðtted with pure 207Pb
(the other, weaker hyperÐne component being hidden by
platinum lines), as discussed in ° 5 of Paper I. A plausible
case can be made that lead is present in the form of pure
207Pb.
Several Si I lines are far too strong if KuruczÏs f-values are
used. To avoid obscuring other spectral features in the synthetic spectrum, we have arbitrarily reduced the transition
probabilities of the lines marked with a double dagger by 1
dex. It remains to be resolved whether this represents a
problem with the transition probabilities, an ionization
anomaly for Si I, or both.
For the ground-state C I line at 1560.309 AŽ , we have
marked the wavelengths at which interstellar absorption
might be expected, but at most one of the two components
seen in the 2343.495 Fe II line (Fig. 29) may be visible here.
Observation Z2MB010M, 1644.73È1653.64 AŽ , Figure
13.ÈThe Ðt for the Ge II line at 1649.194 AŽ is rather poor,
but as this is the strongest germanium line in our observations, we have used it to determine the preliminary abundance used in these synthetic spectra (Paper I, ° 4). The
strongest Hg II resonance line is at 1649.947 AŽ , and an
especially useful Hg III line is at 1647.482 AŽ .
Observation Z0IX010R, 1736.50È1746.27 AŽ , Figure 14.È
Analysis of the Hg III lines near 1738.5 AŽ was discussed by
Leckrone et al. (1993b). The abundance needed to match
these lines is much larger than that needed to match lines of
the dominant ionization stage, Hg II. Further discussion can
be found in Proffitt et al. (1999).
Note the good Ðt to the complex blend between 1741.5
and 1741.7 AŽ , which uses only Kurucz gf values, with
adjustments to some wavelengths made using FTS data.
The abundances used for the elements in this blend were all
determined using optical data (Wahlgren et al. 1994). This is
but one of many examples of where the Kurucz database
and the optical abundance determinations need little or no
improvement.
The Au II line near 1740.47 AŽ contains nine hyperÐne
components over a 0.007 AŽ interval. For clarity, only one
component is marked. Full details can be found in Wahlgren et al. (1995). Note also the Au III line at 1746.047 AŽ .
The N I lines near 1742.72 and 1745.25 AŽ have been used
to set our adopted nitrogen abundance (Paper I, ° 3).
Observation Z0IX010J, 1844.31È1853.44 AŽ , Figure
15.ÈThe ground-state Hg I lines, such as the 1849.508 AŽ
resonance line, are weaker than predicted by LTE models
using the abundance derived from strong Hg II lines. See
Proffitt et al. (1999) for further discussion of this ionization
anomaly.
Observation Z13J010M, 1862.87È1873.10 AŽ , Figure
16.ÈA good example of the indirect level mixing discussed
by Johansson et al. (1995) involves the Fe II lines at 1870.703
and 1870.616 AŽ .
The Hg II line at 1869.226 AŽ is discussed by Proffitt et al.
(1999).
The transition probabilities for the Si II lines at 1868.766,
1869.319, 1870.230, and 1870.784 AŽ are from the calculations of Artru et al. (1981). They are about 2 dex smaller
than the values given in Kurucz (1993), due to the inclusion
of additional conÐguration-interaction e†ects in the Artru
et al. calculations (also see Lanz & Artru 1985, and
Figures 7 and 8 of Leckrone et al. 1990). Use of these transition probabilities reduces, but does not eliminate, the discrepancy between the observed and synthetic spectra.
Observation Z28H010M, 1904.60È1915.57 AŽ , Figure
17.ÈThe two hyperÐne components of a ground-state
FIG. 7.ÈObservation Z2MB010H
FIG. 8.ÈObservation Z3650308
1518
FIG. 9.ÈObservation Z2MB010I
1519
FIG. 10.ÈObservation Z365030B
1520
FIG. 11.ÈObservation Z2MB010L
1521
FIG. 12.ÈObservation Z3650307
1522
FIG. 13.ÈObservation Z2MB010M
1523
FIG. 14.ÈObservation Z0IX010R
1524
FIG. 15.ÈObservation Z0IX010J
1525
FIG. 16.ÈObservation Z13J010M
1526
FIG. 17.ÈObservation Z28H010M
1527
1528
BRANDT ET AL.
intercombination transition of Tl II are found at 1908.572
and 1908.709 AŽ . These lines were discussed extensively by
Leckrone et al. (1996b) and Johansson et al. (1996).
This wavelength interval also contains a number of
examples of ““ superanomalous ÏÏ iron group lines (Leckrone
et al. 1993a), where the line data of Kurucz predicts a strong
line that is weak or completely absent in the observed spectrum. These include Cr II 1904.896, 1904.953, Ti II 1906.239
AŽ , and Fe II 1906.326 AŽ .
Observation Z0G7010J, 1936.74È1947.06 AŽ , Figure 18.È
This observation, from 1991 February, was the Ðrst GHRS
SSA echelle spectrum obtained of s Lupi. The line list for
this interval has been extensively revised, and documenting
all of the changes and corrections made will have to be
deferred until a future paper. Many of the lines were discussed in Leckrone et al. (1993b). The Hg II resonance line
at 1942.287 AŽ was discussed by Leckrone, Wahlgren, &
Johansson (1991), several Ru II lines were analyzed by
Johansson et al. (1994), and the isotope structure of the Pt II
line near 1944.467 was discussed by Kalus (1997), and Kalus
et al. (1998). The discrepancy between the observed and
computed line strengths of the As I line at 1937.594 AŽ arises
because the larger arsenic abundance of Wahlgren et al.
(1994) based on As II lines was used. A similar ionization
anomaly between lines of Zr II (1938.500 AŽ ) and Zr III
(1937.216, 1940.236, 1941.055, and 1946.573 AŽ ) can be seen
in this spectrum (Leckrone et al. 1993b).
Observation Z16C0107, 1997.73È2008.66 AŽ , Figure
19.ÈIt should be remembered that the convention of converting wavelengths above 2000 AŽ to air wavelengths leads
to an ambiguity for wavelengths between 1999.353 and
2000 AŽ . As this observation spans the traditional vacuumair boundary at 2000 AŽ , all wavelengths in this Ðgure are
given as vacuum wavelengths.
In contrast to the observations at shorter wavelengths,
some long stretches of apparently line-free continuum are
visible in this and most of the longer wavelength observations.
Observation Z16C0108, 2012.17È2022.72 AŽ , Figure 20.È
This and subsequent Ðgures use air wavelengths. Note the
Mo II lines at 2015.109 and 2020.314 AŽ (Paper I, ° 4).
Observation Z16C010B, 2023.00È2033.45 AŽ , Figure
21.ÈThe Zn II resonance line at 2025.483 AŽ is extremely
weak or absent, as discussed in Paper I, ° 3. Comparison
with observations that show strong interstellar Fe II absorption (Figs. 29 and 31) suggests that the feature near 2025.3 AŽ
may be partially due to interstellar absorption by this zinc
resonance line.
The transition probabilities given by Kurucz for the
marked P I lines (2023.480, 2024.517, and 2033.474 AŽ ) were
replaced by values from Fawcett (1986) ; however, the
resulting synthetic spectrum still Ðts these lines quite
poorly.
Observation Z16C010C, 2059.71È2071.23 AŽ , Figure
22.ÈThe stellar Zn II resonance line at 2062.004 AŽ would
produce a strong line if present with a solar system abundance. The apparent absence of this line allows an upper
limit to the zinc abundance to be set, as illustrated in Paper
I, ° 3. Some interstellar Zn II absorption may contribute to
the feature near 2061.8 AŽ , and the wavelengths corresponding to the velocity displacements of the ISM Fe II lines seen
in Figure 29 are marked here.
Observation Z16C010F, 2137.18È2149.17 AŽ , Figure 23.È
Wahlgren et al. (1995) used the strong Pt II resonance line
near 2144.25 AŽ to derive a platinum abundance for s Lupi.
We have modeled this line assuming the 60 :40 ratio of
196Pt :198Pt and the isotopic splitting used by Kalus et al.
(1998) but have only marked the average wavelength in the
Ðgure.
Note the resonance line of Cd II at 2144.393 AŽ , from
which we have derived the cadmium abundance (Paper I,
° 4).
The feature near 2148.02 AŽ is signiÐcantly shifted from
the position expected for pure 204Hg, and it is possible that
the Hg II line is blended with a Gd III line for which only
limited atomic data are available (see Paper I, ° 4).
Observation Z16C010G, 2147.73È2159.53 AŽ , Figure
24.ÈWe do not have a reliable gf value for the Ir II line at
2152.708 AŽ and have not yet attempted to derive an iridium
abundance.
A derivation of an upper limit to the tin abundance from
the 2151.518 AŽ Sn II intercombination line is discussion in
Paper I, ° 4.
The Ðts to the P I lines in this spectrum (2149.142,
2152.939, 2154.072, and 2154.113 AŽ ) are quite poor.
Observation Z16C010J, 2201.25È2212.05 AŽ , Figure
25.ÈA very large blemish a†ects data between 2202.3 and
2202.85 AŽ . While we have removed the data a†ected by this
blemish, the S/N of the remaining data in this region is
signiÐcantly reduced.
The Pt II isotope structure for the 2202.02 AŽ line is from
Kalus et al. (1998).
The Pb II line at 2203.533 AŽ is weaker but less blended
than the resonance line at 1433.905 AŽ and could be used to
set a more secure but less restrictive upper limit to the lead
abundance. However, there is clearly missing opacity near
this wavelength, which would need to be accurately
modeled.
No trace of the W II 2204.482 AŽ line is visible, which
allows an upper limit to the tungsten abundance to be set
(Wahlgren et al. 1998).
Note the Pd II line at 2202.354 and 2207.484 AŽ (Lundberg
et al. 1996).
Observation Z16C010K, 2262.83È2274.57 AŽ , Figure 26.È
Noteworthy in this interval are the Al I line at 2263.463 AŽ
and the Cd II line at 2265.019 AŽ (Wahlgren et al. 1995).
Observation Z16C010N, 2271.39È2282.97 AŽ , Figure 27.È
There are numerous isotopic and hyperÐne components of
Re II between 2275.14 and 2275.34 AŽ , but there is no trace of
them in the observed spectrum. Only the shortest and
longest wavelength components of this blend are marked. A
more detailed discussion can be found in Wahlgren et al.
(1997).
Observation Z16C010O, 2323.93È2336.75 AŽ , Figure 28.È
Another example of indirect level mixing can be found in
the Fe II lines at 2325.380 and 2325.246 AŽ (Johansson et al.
1995).
Note the Pd II line at 2336.587 AŽ used for abundance
determination by Lundberg et al. (1996).
Observation Z16C010R, 2335.20È2347.84 AŽ , Figure
29.ÈThe feature near 2343.3 AŽ is due to interstellar absorption from Fe II 2343.495 AŽ . For reference, we have marked
the location of this line for shifts of [6 and [12 km s~1
from the center-of-mass velocity of the binary.
The Ba II lines at 2335.270 and 2347.594 AŽ are discussed
in Paper I, ° 4.
Observation Z0IX010O, 2346.95È2359.37 AŽ , Figure
30.ÈThe absence of the As I line predicted at 2349.839 AŽ is
FIG. 18.ÈObservation Z0G7010J
FIG. 19.ÈObservation Z16C0107
1530
FIG. 20.ÈObservation Z16C0108
1531
FIG. 21.ÈObservation Z16C010B
1532
FIG. 22.ÈObservation Z16C010C
1533
FIG. 23.ÈObservation Z16C010F
1534
FIG. 24.ÈObservation Z16C010G
1535
FIG. 25.ÈObservation Z16C010J
1536
FIG. 26.ÈObservation Z16C010K
1537
FIG. 27.ÈObservation Z16C010N
1538
FIG. 28.ÈObservation Z16C010O
1539
FIG. 29.ÈObservation Z16C010R
1540
FIG. 30.ÈObservation Z0IX010O
1541
FIG. 31.ÈObservation Z16C010S
1542
FIG. 32.ÈObservation Z16C010V
1543
FIG. 33.ÈObservation Z16C010W
1544
FIG. 34.ÈObservation Z0IX010G
1545
FIG. 35.ÈObservation Z16C010X
1546
FIG. 36.ÈObservation Z28H010L
1547
1548
BRANDT ET AL.
due to a combination of the ionization anomaly relative to
As II (see Paper I, ° 4), as well as to the poor quality of the
adopted transition probability for this line. Also noteworthy in this interval are a number of strong, lowexcitation Fe II lines (2348.117, 2348.303, and 2359.106 AŽ ),
and two Pd II lines (2351.347 and 2357.633 AŽ ).
Observation Z16C010S, 2376.15È2387.90 AŽ , Figure
31.ÈOne subexposure of this observation was eliminated
due to inconsistent line proÐles. Note the Pd II line at
2377.923 AŽ (Lundberg et al. 1996).
The Fe II lines at 2379.171 and 2379.311 AŽ show the
e†ects of strong indirect level mixing (Johansson et al. 1995),
although in this case the lines are badly blended with other
Fe II lines.
The Ñat-bottomed feature centered near 2381.85 AŽ is
interstellar Fe II at 2382.038 AŽ . We have marked the location of this line for shifts of [6 and [12 km s~1 from the
center-of-mass velocity of the binary.
Observation Z16C010V , 2407.62È2421.18 AŽ , Figure
32.ÈThe Pt II line near 2420.82 AŽ is one of the lines
analyzed by Kalus et al. (1998). The lines of Rh II at
2415.840 and 2420.968 AŽ were used to derive the abundance
of rhodium (Lundberg et al. 1998).
Observation Z16C010W , 2434.82È2448.04 AŽ , Figure 33.È
Wavelengths for two Fe II lines (2436.418 and 2436.565 AŽ )
showing the e†ects of strong indirect level mixing are taken
from Johansson et al. (1995). See footnote a of Table 1 in
Paper I. Note the Pd II line at 2435.321 AŽ (Lundberg et al.
1996).
Observation Z0IX010G, 2528.21È2542.33 AŽ , Figure
34.ÈThe Hg I ground-state intercombination line at
2536.539 AŽ shows a strong ionization anomaly relative to
the abundance determined from the majority ionization
state, Hg II (see Proffitt et al. 1999).
Observation Z16C010X, 2602.06È2614.81 AŽ , Figure
35.ÈThe f-value from Kurucz for the Mn II resonance line
at 2605.684 was replaced by a larger transition probability
taken from Morton (1991), but the predicted line is still
much too weak. We have not investigated whether substantial hyperÐne structure is expected for this line.
The wings of two low-lying Fe II lines at 2607.088 and
2611.874 AŽ are poorly Ðtted. See footnote b of Paper I,
Table 1.
Observation Z28H010L , 2673.98È2688.32 AŽ , Figure
36.ÈOne of the four subexposures of observation
Z13J010LM shows inconsistent line proÐles, and this subexposure was not used.
Note the Mo II lines at 2683.229 and 2684.140 AŽ (Paper I,
° 4). The zirconium ionization anomaly is again illustrated
by the comparison of the Zr II line at 2678.646 AŽ with the
Zr III line at 2682.181 AŽ (Paper I, ° 4).
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