1. introduction - Lund Observatory

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.

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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

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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

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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

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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

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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/.

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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

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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

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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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