MERCURY IN THE HgMn STARS s LUPI AND HR 7775 CHARLES R. PROFFITT1,2 Science Programs, Computer Sciences Corporation; hrspro itt=hrs.gsfc.nasa.

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1 THE ASTROPHYSICAL JOURNAL, 512:942È960, 1999 February 20 ( The American Astronomical Society. All rights reserved. Printed in U.S.A. MERCURY IN THE HgMn STARS s LUPI AND HR 7775 CHARLES R. PROFFITT1,2 Science Programs, Computer Sciences Corporation; hrspro itt=hrs.gsfc.nasa.gov TOMAS BRAGE Department of Physics, University of Lund, Box 118, S Lund, Sweden; brage=kurslab.fysik.lu.se DAVID S. LECKRONE3 Laboratory for Astronomy and Solar Physics, NASA Goddard Space Flight Center, Greenbelt, MD 20771; hrsleckrone=hrs.gsfc.nasa.gov GLENN M. WAHLGREN1,4 Science Programs, Computer Sciences Corporation; spek glenn=garbo.lucas.lu.se JOHN C. BRANDT3 Laboratory for Atmospheric and Space Physics, Campus Box 392, University of Colorado; brandt=lyrae.colorado.edu CRAIG J. SANSONETTI AND JOSEPH READER National Institute of Standards and Technology, Gaithersburg, MD 20899; Craig.Sansonetti=nist.gov, Joseph.Reader=nist.gov AND SVENERIC G. JOHANSSON5 Department of Physics, University of Lund, Box 118, S Lund, Sweden; spek sej=garbo.lucas.lu.se Received 1997 December 15; accepted 1998 September 24 ABSTRACT Observations of mercury lines in the HgMn stars s Lupi and HR 7775 made with the Hubble Space T elescope Goddard High Resolution Spectrograph are presented and analyzed. In s Lupi we Ðnd that all observed lines are consistent with the same isotopic mixture (essentially pure 204Hg). Strong ionization anomalies are present, with UV Hg I lines being too weak and Hg III lines too strong for the abundance derived from lines of the majority ionization state, Hg II. Observations of mercury in HR 7775 show less extreme isotope and ionization anomalies. We Ðnd that the ionization anomaly in the Hg I resonance lines can be plausibly explained as a non-lte e ect, but the same non-lte calculations show that the Hg III ionization anomaly in s Lupi cannot be explained in this way. Radiative force calculations show that the observed mercury abundance cannot be supported in the atmosphere by the radiative forces alone. We suggest that weak mixing brings mercury into the line-forming region from below the photosphere, while a wind of order 10~14 M yr~1 supports a cloud of Hg III at very small optical depths. _ Subject headings: di usion È stars: abundances È stars: individual (s Lupi, HR 7775) È stars: peculiar 1. INTRODUCTION Among the most puzzling of the abundance peculiarities present in chemically peculiar stars are those of mercury, platinum, gold, and thallium. In the cool HgMn star s Lupi, each of these elements appears to be overabundant relative to solar system abundances by 4È5 orders of magnitude, and for those with more than one stable isotope (mercury, platinum, and thallium), the mixture of each is much more strongly weighted toward the heavier isotopes than in the solar system mixtures (Anders & Grevesse 1989). In addition, very large ionization anomalies are found for these elements, with dramatically di erent abundances being derived from lines of di erent ionization states. For mercury in s Lupi, White et al. (1976) found that about 99% was in the form of the heaviest isotope, 204Hg, 1 Institute for Astrophysics and Computational Science, Catholic University of America; and Goddard High Resolution Spectrograph Science Team. 2 Laboratory for Astronomy and Solar Physics, NASA Goddard Space Flight Center, Code 680, Greenbelt, MD 20771; and Space Telescope Science Institute. 3 Goddard High Resolution Spectrograph Investigation DeÐnition Team. 4 Department of Physics, University of Lund, Box 118, S Lund, Sweden. 5 Lund Observatory, University of Lund, Box 43, S Lund, Sweden. 942 which is only about 7% of the solar system mixture. Kalus et al. (1998) found that platinum in s Lupi is dominated by a roughly equal mixture of the two heaviest isotopes, 196Pt and 198Pt, which respectively comprise only 25% and 7% of the platinum in the solar system. Leckrone et al. (1996) found a similar overabundance of 205Tl relative to 203Tl. Similar, though usually less extreme, isotope anomalies are found in most other cool (T \ 12,500 K) HgMn stars, including HR 7775 (Smith 1997; eff Kalus et al. 1998). Since no plausible nucleosynthesis scenario has been proposed to explain the overabundances and isotope anomalies of these heavy elements, it is usually assumed they can be explained by radiatively driven atomic di usion. Scenarios of this sort were outlined by Michaud, Reeves, & Charland (1974). In their picture, the strong radiative forces on Hg II would push the mercury upward until it reaches the small optical depth where the ionization equilibrium favors Hg III. The much smaller radiative force on the lines of Hg III would lead to the accumulation of a dense cloud of mercury at small optical depths. As the abundance in this cloud increases, the saturation of the Hg II lines would lead to an increasingly large region where the radiative and gravitational forces on mercury approximately cancel, and the small mass di erence between the isotopes would then lead to their segregation. The lighter isotopes would either be hidden at small optical depths in the form of Hg III or preferentially removed from the stellar photosphere in a

2 MERCURY IN s LUPI AND HR wind. The lack of detailed atomic data for mercury prevented Michaud et al. (1974) from presenting more than an outline of these possible mechanisms. Mercury also provides what is perhaps the most extreme example of the ionization anomalies observed in s Lupi. As was Ðrst shown by Leckrone et al. (1993), Ðtting the A Hg III line with a standard LTE model requires a mercury abundance D15È20 times larger than is required to Ðt the Hg II resonance line at 1942 A (Leckrone, Wahlgren, & Johansson 1991). Leckrone et al. (1993) also demonstrated that the Hg III line shows the same isotope mixture as the Hg II lines, showing no indication of the missing ÏÏ lighter isotopes. Over the past several years, the Hubble Space T elescope Goddard High-Resolution Spectrograph (GHRS) has obtained a number of additional observations with extremely high resolutions (3È4 kms~1) and high signal-tonoise ratios (S/N up to 100:1) that cover mercury lines in the cool HgMn stars s Lupi and HR We will take advantage of these new observations by using modern atomic data and model atmosphere codes to carry out a more complete and accurate set of computations than has previously been possible, with the ultimate goal of addressing the di usion paradigm in as comprehensive a manner as possible. In this paper we will attempt to see whether or not the mercury ionization anomalies observed in the GHRS data might be explainable purely in terms of non-lte e ects, without invoking abundance stratiðcations in the atmosphere. We will also make more realistic estimates of the radiative forces on mercury than have previously been possible, to determine whether radiative forces alone are capable of explaining the observed abundance anomalies. We have previously performed similar radiative force and non-lte calculations for thallium in s Lupi (Leckrone et al. 1996; Proffitt, Brage, & Leckrone 1996). We concluded that the radiative force on thallium in the atmosphere of s Lupi could not, by itself, support sufficient thallium at small optical depths to explain the observed line proðles. There is sufficient radiative force to support the observed thallium abundance just below the photosphere, and this suggests that the atmosphere is mixed. Mercury will provide an even better test of the di usion hypotheses than thallium, since there is a larger set of observed mercury lines, covering a wider range of ionization states and excitation energies Previous Mercury Abundance Determinations The Hg II line near 3984 A is the lowest excitation line of Hg II observable with ground-based instruments, and it also has a very large isotope shift (0.23 A between 204Hg and 198Hg; see 2.2). Thus, the lines of the even isotopes can be easily separated at resolutions of 105 or more. While each of the odd isotopes has several hyperðne components, they are not easily separated from the lines of the even isotopes in stellar spectra. One disadvantage of this line is that at the abundances found in HgMn stars, the individual components are on the saturated part of the curve of growth. This makes derived abundances and isotope ratios sensitive to the adopted broadening parameters and microturbulence. Until recently, the most detailed analysis of the 3984 A line in HgMn stars was that of White et al. (1976). While they reported the observed central wavelength of this feature for 30 stars, for only three of the most slowly rotating ones ( CrB, s Lupi, and HR 4072) do they have data that allow the abundance of individual mercury isotopes to be estimated. In s Lupi only 204Hg and 202Hg could be directly measured, and the heavier isotope was found to be 100 times more abundant than the lighter. In HR 4072, 204Hg, 202Hg, and the blended 200Hg ]201Hg feature were measured, and in CrB, the 198Hg ]199Hg blend could be measured in addition to the heavier isotopes. From these observed abundances, they suggested that the relative isotope mixture in a given star could be Ðtted with the formula [N(A)/N(202)] * \ exp [q(a [ 202)], [N(A)/N(202)] _ where N(A) is the number abundance of the isotope with atomic number A, and q is a dimensionless mix parameter that di ers from star to star. Table 1 illustrates how the isotope mixture varies as a function of q. The solar system mixture (Anders & Grevesse 1989) is obtained for q \ 0. Assuming that this formula describes the isotope mix, White et al. (1976) then derived q-values for the rest of their observations using the measured central wavelengths of the Hg II 3984 A feature. For HgMn stars with T [ 12,500 K, eff they found that q ¹ 1.0, with a majority of stars apparently having an isotope mixture similar to that in the solar system. In contrast, most HgMn stars with T \ 12,500 K eff show marked isotope anomalies, and have 1 \ q \ 1.5, with a few having larger q (including q \ 3 for s Lupi), but apparently none have a solar system mixture. Smith (1997) recently discussed new observations of the 3984 A line for a number of HgMn and chemically normal late-b stars, and applied modern spectral synthesis techniques to the analysis. He found a distribution of abundances and isotope anomalies similar to that found by White et al. (1976). However, for the star HR 7775 he concluded that the q-parameterization could not model the isotope mixture, instead Ðnding a 202Hg/204Hg ratio in HR 7775 of about 1:2, which is best Ðtted by q B 1.3, while the ratio of the lighter isotopes (especially 200Hg and 201Hg) to 204Hg was better Ðtted by q B 1.8. M. M. Dworetsky (1998, private communication), Wahlgren (1999), S. Hubrig, F. Castelli, & G. Mathys (1998, private communication), and V. M. Woolf (1998, private communication) also report TABLE 1 FRACTIONAL ABUNDANCE OF EACH MERCURY ISOTOPE AS A FUNCTION OF q q

3 944 PROFFITT ET AL. Vol. 512 similar results for the isotope mixture in this star. From this line, Smith also derived a mercury abundance of log N(Hg) \ 6.40 ^ 0.45 for s Lupi and 6.33 ^ 0.36 for HR 7775 [on the scale where log N(H) \ 12]. Smith (1997) also considered IUE observations of the 1942 A line, and found log N(Hg) \ 6.20 ^ 0.20 for s Lupi and 6.30 ^ 0.20 for HR It should also be remembered that with the 3984 A line, neither of the odd isotopes can be measured separately, since they are severely blended with lines of the even isotopes. Since the lines of the odd isotopes show substantial hyperðne structure, it is possible that an isotopic separation mechanism that depends sensitively on the details of the line spectrum could result in di erent abundance patterns for the odd and even isotopes. A direct measurement of the abundance of either 199Hg or 201Hg would therefore provide useful theoretical constraints. In the cool HgMn stars, subordinate lines of Hg I can be observed in optical spectra. Abundances have been reported using the 4358 A Hg I line by Adelman (1994), who found log N(Hg) \ 6.1 for HR 7775, Wahlgren, Adelman, & Robinson (1994), who found log N(Hg) \ 5.89 for s Lupi, and Smith (1997), who found log N(Hg) \ 6.14 ^ 0.16 and 6.32 ^ 0.16 for s Lupi and HR 7775, respectively. These results are consistent with the abundances found for these stars from the Hg II lines. 2. ATOMIC CALCULATIONS The atomic data needed for this paper serve a number of di erent purposes. First, abundances will be determined by directly comparing observed data with synthetic spectra. Next, non-lte ionization equilibrium models will be calculated to investigate ionization anomalies. Since HgMn stars show isotope anomalies, we will also need to calculate isotope shifts and hyperðne structure for those transitions for which sufficient laboratory data are not available. Finally, we will compute radiative forces on mercury for comparison with the expectations of radiative di usion theory. When comparing a synthetic to an observed spectrum to determine an abundance, available data from the GHRS are in principle of sufficiently high S/N to reduce uncertainties in the abundance ratios to the 10% level, but only if the oscillator strength of the speciðc transition is known to comparable accuracy. This is a demanding requirement; it is difficult to obtain such accuracy with experimental methods, and theoretical approaches in an atom as heavy as mercury are complicated by the need to take both relativistic and core-polarization e ects into account. When constructing a non-lte model, radiative transition probabilities and collisional cross sections are needed for a large number of transitions. For many of these transitions, especially for those between highly excited states, reproducing the collective behavior is more important than achieving extremely high accuracy for individual transitions. This is because collisional processes will keep the relative populations of the closely spaced upper levels close to LTE ratios throughout most of the atmosphere. Indeed, in complex non-lte calculations such levels are often grouped together to reduce the computational burden. This suggests that calculations and experimental data of the highest possible accuracy must be used for the transitions to which non-lte e ects are most sensitive, as well as for transitions that are directly observable, while for the other transitions data of considerably lower accuracy may suffice. To this end, we adopted a mixture of di erent strategies for computing radiative transition probabilities and the other atomic parameters needed. In this paper we will only summarize these calculations. Full details of the techniques used are presented in a companion paper (Brage, Proffitt, & Leckrone 1999) T ransition Probability Calculations For heavy ions such as mercury, a methodology has been developed and tested for large-scale systematic calculations, based on the multiconðguration Dirac-Fock (MCDF) method (Grant et al. 1980) using the GRASP94 package (Parpia, Froese Fischer, & Grant 1996). To compute the most important transition probabilities, we use large-scale restricted active space (MDCF/RAS) calculations (Brage, Leckrone, & Froese Fischer 1996). This method is highly systematic and allows for convergence and accuracy tests. The method focuses on a few lines and includes or probes for the e ects of polarization, or core-valence correlation (CV), of deep core subshells. Our goal is to calculate transition strengths and hyperðne structure constants to better than 10% precision. In order to include larger numbers of levels for which we want calculations of good accuracy, we use the crosswise optimization (MCDF-CO) method described by Brage et al. (1993). This is designed to be more Ñexible and to simultaneously treat a larger number of levels than the MCDF- RAS method, but it only includes the most important CV e ects. For the strongest transitions, our goal with these techniques is a precision of 15% or better. To complete our model ions, we also perform semiempirical calculations (Cowan 1981) for a large number of levels. While the Cowan code calculations necessarily include a less accurate treatment of core-polarization e ects than do the MCDF calculations, they do allow the inclusion of a much larger number of levels. They may therefore include level-mixing e ects missed by the more sophisticated but more limited MCDF calculations. The MCDF calculations are generally expected to give superior results for transitions between low-lying levels, but in some other cases the Cowan code results may be preferred. The MCDF-RAS calculations were only performed for the two resonance lines of Hg II near 1942 and 1649 A. Transition probabilities were calculated for 38 other transitions of Hg II using the MCDF-CO method and for 155 transitions using the Cowan code. For Hg III, 36 transitions were treated using MCDF-CO, and 221 using the Cowan code. For Hg I we calculated 1063 transition probabilities with the Cowan code, but substituted an average of theoretical and experimental values from the literature for the ground-state resonance lines at 2536 and 1849 A (cf. Migdalek & Stanek 1980). The Ðnal adopted transition probabilities for lines discussed in this paper are summarized in Table 2. The wavelengths and energy levels given in this table are our adopted values for 204Hg (see 2.2). Radiative and Stark damping constants based on our model ions are also listed Isotope Shifts and HyperÐne Splitting. While the size of the isotope shift varies considerably from one atomic energy level to the next, apart from an overall scale factor, the relative spacings of the di erent even isotopes are the same for all energy levels and lines,

4 No. 2, 1999 MERCURY IN s LUPI AND HR TABLE 2 PARAMETERS FOR SELECTED LINES OF 204Hg j 204 E lo E hi (A ) Ion log gf (cm~1) J lo (cm~1) J hi log! R log (! S /n e ) Hg I [ Hg I [ [ Hg II [ Hg II [ Hg II [ [ Hg II [ Hg II [ [ Hg II [ Hg II [ [ Hg II [ [ Hg III [ Hg III [ [ Hg III [ [6.45 (e.g., for any line or level the 204Hg-200Hg shift is times the 202Hg-200Hg shift). This is only approximately true for the odd isotopes because of second-order hyperðne interactions, but that e ect should be small for the levels considered in this paper. We have adopted the values of the relative shifts as given by Sansonetti & Reader (1993). Wavelengths for Hg II and III lines for natural mercury have been measured by Sansonetti & Reader (1999). For a number of Hg II levels, isotope shifts can be estimated by comparing these results to those of Reader & Sansonetti (1986) for lines of 198Hg II. Sansonetti & Reader (1993) measured isotope shifts for several Hg III lines. For all levels included in our MCDF calculations, isotope shifts were also calculated using the techniques described in Brage et al. (1999). In the neutral mercury atom and its next two ionization stages, the 6s states interact more strongly with the nuclear potential than do other optically active states, and the TABLE 3 ADOPTED ISOTOPE AND HYPERFINE STRUCTURE j (A ) (0.333) 0.495(0.667) (0.500) 0.502(0.167) 0.500(0.333) (0.667) 0.539(0.333) (0.167) 0.538(0.500) 0.508(0.333) (0.028) 0.708(0.556) (0.040) 0.713(0.120) 0.713(0.040) (0.389) 0.709(0.028) 0.712(0.335) 0.713(0.087) 0.713(0.037) 0.708(0.556) 0.712(0.201) 0.713(0.051) 0.713(0.037) 0.712(0.051) (0.562) 0.739(0.021) (0.125) 0.716(0.001) 0.740(0.030) (0.417) 0.701(0.030) 0.718(0.039) 0.741(0.251) 0.701(0.001) 0.719(0.179) 0.769(0.344) (0.250) 0.941(0.125) (0.031) 0.928(0.156) 0.954(0.156) (0.625) 0.928(0.438) 0.953(0.063) 0.955(0.156) (0.250) 0.226(0.125) (0.031) 0.223(0.156) 0.229(0.156) (0.625) 0.223(0.438) 0.229(0.063) 0.229(0.156) (0.250) 0.284(0.250) (0.313) 0.293(0.063) 0.262(0.313) (0.500) 0.300(0.313) (0.375) 0.710(0.042) (0.375) 0.710(0.233) 0.712(0.131) (0.583) 0.710(0.058) 0.710(0.073) 0.712(0.056) 0.710(0.004) 0.710(0.006) 0.713(0.063) (0.313) 0.780(0.562) (0.350) 0.781(0.100) 0.783(0.063) (0.063) 0.781(0.063) 0.781(0.087) 0.781(0.087) 0.783(0.025) 0.781(0.125) 0.783(0.100) 0.784(0.063) (0.400) 0.506(0.029) (0.002) 0.500(0.003) 0.507(0.040) (0.571) 0.492(0.041) 0.501(0.052) 0.508(0.160) 0.494(0.357) 0.502(0.245) 0.511(0.100) (0.400) 0.476(0.571) (0.357) 0.481(0.052) 0.488(0.040) (0.029) 0.471(0.245) 0.481(0.002) 0.488(0.003) 0.471(0.041) 0.482(0.160) 0.489(0.100) (0.560) 0.528(0.360) (0.050) 0.523(0.050) 0.528(0.070) (0.040) 0.529(0.040) 0.519(0.050) 0.527(0.173) 0.534(0.343) 0.522(0.080) 0.527(0.057) 0.534(0.131) 0.522(0.070) NOTE.ÈThe fractional part of the wavelength listed under each isotope should be added to the integer wavelength in the Ðrst column. For the odd isotopes, the fraction of the total gf-value to be assigned to each component is given in parentheses.

5 946 PROFFITT ET AL. Vol. 512 isotope shifts are therefore largest for lines between states with di erent numbers of 6s electrons. When the lower level of the transition has the larger number of 6s electrons (the usual case), the line is at a longer wavelength in the heavier isotope. For transitions in Hg II that change the number of 6s electrons by 1, the wavenumber typically shifts by about 0.2È0.3 cm~1 between 202Hg and 200Hg, while lines between the 5d96s2 levels and those with no 6s electrons are shifted by about 0.5 cm~1. Transitions that involve no change in the number of 6s electrons have very small isotope shifts. For a few levels where we have not calculated the isotope shift explicitly, we have adopted the average shift from levels of similar conðguration. The odd isotopes of mercury (199Hg and 201Hg) have a nonzero nuclear spin, and most lines of mercury therefore show hyperðne structure in addition to the isotope shift. The hyperðne splittings vary much less predictably from level to level than do the isotope shifts, and for a particular line this splitting may be either larger or smaller than the isotope shifts. For the levels included in our MCDF calculations, we have calculated the hyperðne A and B parameters (see Brage et al. 1999). Our calculated hyperðne parameters di er by less than 3% from those measured by Guern et al. (1977) from the 1942 A resonance line. For those lines for which we have calculated the structure parameters for both levels, we have used these data to predict the wavelengths for each component of each isotope. Where we have not calculated the hyperðne parameters for each level, we have simply included in our line list a single component at the unperturbed wavelength for that isotope. For the ground-state and 2P levels of Hg II, we 1@2 have used the experimental values for the hyperðne parameters in place of our own calculations. For Hg I, energy levels and isotope shifts were taken from Gerstenkron, Labarthe, & Vergès (1977) and Kaufman (1962), while the hyperðne structure for 6sÈ6p transitions of the odd isotopes are from Lecler (1968), Kohler (1961), and Stager (1963). In this ion, we have only included the hyper- Ðne structure for these two resonance lines. Line data for selected transitions are listed in Table 2, and the adopted isotope and hyperðne splittings for these lines are given in Table 3. The radiative widths given in Table 2 are calculated using our transition probabilities. The Stark broadening parameters were calculated using the modiðed semiempirical approximation discussed by Gonzalez et al. (1995) and are given at a temperature of 10,000 K. For the highest levels of Hg II we consider (E º 11,000 cm~1), the incompleteness of our model ion has probably led us to underestimate the Stark widths. Nowhere in our model atmospheres is less than 40% of the hydrogen ionized, implying that Stark broadening will always dominate van der Waals broadening. Complete tables of our adopted parameters for the lines of each isotope will be made available on the web pages of the SAM (Systematic, Accurate, MulticonÐguration calculations) project.6 The wavelengths given in Table 3 for Hg II and Hg III were determined by applying the calculated relative isotope shifts to our best estimate of the mean energy levels for natural ÏÏ mercury. These values may di er slightly from published laboratory values, and the data in Table 3 should not be considered a substitute for precision laboratory measurements. However, errors for most lines in Table 3 should be less than 2 ma, and this is sufficient for our purposes. 3. LTE ANALYSIS OF GHRS OBSERVATIONS The GHRS observations containing mercury lines discussed in this paper are listed in Tables 4 and 5. We consider only echelle exposures taken through the small science aperture (SSA). For s Lupi, we have GHRS data for a number of relatively unblended mercury lines, including two or more lines for each of the Ðrst three ionization stages. For HR 7775, our UV data set is much more limited, and includes only the 1942 A Hg II resonance line and the A Hg III line. The observations were obtained as part of the GHRS GTO programs of Brandt and Leckrone, or as part of the STScI science veriðcation and calibration programs. Most of the observations listed in Table 4 and all of those listed in Table 5 used the FP-SPLIT \ 4 option to reduce Ðxedpattern noise. However, the two observations of s Lupi that list Schultz as the PI were made as ACCUM exposures at a single carousel position. These data must be used with caution, since without the FP-SPLIT option there is no reliable way to detect or eliminate Ðxed-pattern noise or blemishes, and this can lead to sizable artifacts in the spectrum. To allow comparisons of the relative quality of the observations, we list the maximum value of the signal-to-noise 6 Data available online at: TABLE 4 GHRS SSA ECHELLE OBSERVATIONS OF s LUPI j min j max Start Time t exp j A [ j B (A ) (A ) Root Name RJDa (days) PI (s) S/N (A ) Comments Z2MB010B Leckrone Tl II, HgII Z2MB010D Leckrone Hg III Z2MB010M Leckrone Hg II/III Z0IX010R Leckrone [0.451 Hg III, miscentered Z0IX010J Leckrone [0.442 Hg I, miscentered Z13J010M Gilliand [0.743 Hg II Z13J510M Gilliand Z0G7010J Ebbets [0.799 Hg II Z0IX010B Leckrone [0.426 miscentered Z3D5020C Schultz Hg II, ACCUM Z3D5020U Schultz Hg II, ACCUM Z0IX010G Leckrone [0.585 Hg I, miscentered a Reduced Julian day is the Julian day [ 2,400,000.

6 No. 2, 1999 MERCURY IN s LUPI AND HR TABLE 5 GHRS SSA ECHELLE OBSERVATIONS OF HR 7775 j min j max Start Time t exp (A ) (A ) Root Name RJD (days) PI (s) S/N Comments Z Brandt Hg III Z365010A Brandt Hg II ratio (S/N) per substep bin for each observation, as calculated using Poisson statistics. Each substep bin corresponds to a quarter of a diode on the GHRS detector array, or about one-half of the instrumental point-spread function. Because of the Ðxed-pattern noise in the GHRS, which is only partially ameliorated even with the use of the FP- SPLIT option, the true continuum S/N ratio will be less. For s Lupi, the wavelength o set between the spectra of the primary and secondary stars was calculated using the orbital parameters determined by Dworetsky (1972). Additional details of our data reduction and analysis procedures can be found in Wahlgren et al. (1995). Most of the observations of s Lupi listed in Table 4 will be published in full in a spectral atlas of s Lupi GHRS observations (Brandt et al. 1998). Any reader who wishes to examine the data for lines discussed but not plotted in this paper, or who wishes to examine a larger portion of a spectrum than is shown here, may Ðnd this fuller account useful. A similar publication of the more limited GHRS observational data for HR 7775 will follow at a later date. In this section, our analysis of the mercury lines detected in GHRS observations of s Lupi and HR 7775 will be made assuming LTE line formation and a uniform chemical composition throughout the atmosphere. While it is possible that these assumptions are inadequate, such an analysis does provide a starting point for more complex models. Except where otherwise stated, we will use the same parameters for s Lupi as those adopted by Wahlgren et al. (1994) in their optical study of this star (see Table 6). The model atmosphere calculated for that paper will be used together with the SYNTHE code and line lists of Kurucz (1993) to compute synthetic spectra. The relative contribution of the secondary in s Lupi will be estimated using the relative Ñuxes (per unit area of the stellar surface) from the synthetic spectra calculations, scaled assuming that the radius of the primary is 1.67 times that of the secondary (Wahlgren et al. 1994) and shifted according to the calculated orbital velocity di erence (Dworetsky 1972). A model atmosphere similar to that for s Lupi has been prepared for HR 7775 (Wahlgren et al. 1999). There is no evidence that more than one star contributes to the Ñux observed for HR The LTE abundance results derived in this section are summarized in Table 7. The abundances in this table and in the discussion below are given as log N(Hg) on a scale where log N(H) \ 12. Expressed in this way, the solar TABLE 6 ADOPTED STELLAR PARAMETERS T eff v sin i m Star Name (K) log g (km s~1) (km s~1) s Lupi A s Lupi B HR system mercury abundance is log N(Hg) \ 1.09 (Anders & Grevesse 1989). The uncertainties quoted for each line are estimates of the uncertainty in Ðtting the synthetic spectra to the observed data, and include an estimate of any uncertainty in setting the continuum level. They are not statistically rigorous error bars, and do not include errors associated with uncertainties in the transition probabilities or with any possible systematic errors in the underlying assumptions Hg II Hg II is the dominant ionization state of mercury in late B and early A stars. Throughout most of the atmosphere of s Lupi, the majority of the mercury atoms should be in the ground state of this ion. While the cores of the strong Hg II resonance lines will be formed at small optical depths and may be subject to substantial non-lte e ects, the wings of these lines will be dominated by radiation damping and formed much deeper in the atmosphere. It is hoped that these line wings will be formed under LTE conditions and give a good indication of the abundance in the bulk of the atmosphere. In addition, since the damping wings of these strong lines are much broader than the Doppler widths, a poorly modeled or unknown blending line of another element would have to be extremely strong to a ect the Ðt in any signiðcant way A (6s 2S È6p 2P ) 1@2 1@ A in s Lupi The Ðrst GHRS observation of the 1942 A resonance line of Hg II (Z0G7010J) was previously discussed by Leckrone et al. (1991). The second observation (Z0IX010B) is entirely consistent with the Ðrst, but because of the higher S/N ratio of the Ðrst observation, we will use Z0G7010J for our TABLE 7 LTE MERCURY ABUNDANCES IN s LUPI AND HR 7775 j 204 (A ) Ion s Lupi HR Hg I 5.2 ^ Hg I 5.1 ^ Hg II 5.7 ^ Hg II 5.3 ^ Hg II 6.1 ^ Hg II 6.4 ^ Hg II 6.24 ^ ^ Hg II 5.8 ^ Hg II 6.2 ^ Hg III 6.5 ^ Hg III 6.8 ^ Hg III 7.4 ^ ^ 0.3 NOTE.ÈThe quoted uncertainties (1 p) include only the uncertainty in Ðtting the synthetic spectrum to the observed data and do not include uncertainties in transition probabilities or possible systematic uncertainties.

7 948 PROFFITT ET AL. Vol. 512 analysis. We do not coadd the two spectra, because the velocity di erence between the primary and secondary di ers between the two observations. Our adopted f-value is 0.05 dex larger than that used by Leckrone et al. (1991), and we have chosen a slightly di erent wavelength alignment and continuum normalization. Since our Ðt to this line is dominated by the damping wings, these changes result in a mercury abundance of log N(Hg) \ 6.24 or [Hg/ H] \]5.19 that is 0.15 dex larger than that found by Leckrone et al. (1991). Our adopted wavelength scale is based on the Ru II lines at and A (Johansson et al. 1994) and the Fe II line at A, for which accurate laboratory wavelengths measured using a Fourier transform spectrometer (FTS) are available. The Ðt to the data with our default parameters is shown in Figure 1. As in Leckrone et al. (1991), the observed line core is noticeably deeper and broader than in the model spectrum. This disagreement could be alleviated to some extent by reducing either the assumed contribution of the secondary or the correction for scattered light. Near 1942 A, the default calibration for the scattered light reduces the Ñux by 3.7% of the mean observed value (the d-coefficient of Cardelli, Ebbets, & Savage 1993), although Cardelli et al. also found that actual measurements of the scattered light vary from their adopted relation by ^1% of the mean Ñux. As an example of how such uncertainties can a ect the Ðtting of the mercury line, we show in Figure 2 how changing the primary-to-secondary continuum light ratio from our default of 6.64:1 to 7.71:1 can improve the Ðt. Despite the improved Ðt, it would be premature to accept this as the proper solution. We cannot exclude the possibility that the discrepancy seen in Figure 1 might be the result of non-lte e ects, partial redistribution of photons within the line proðle, an error in the temperature proðle for the outer layers of the model atmosphere, or vertical stratiðcation of the Hg/H ratio. It is also impossible to exclude the possi- FIG. 2.ÈAs in Fig. 1, except that the contribution of the secondary in the synthetic spectrum calculation has been reduced to improve the Ðt to the line core. A similar result could be obtained by reducing the correction for scattered light. bility that unknown blends might be a ecting the line proðle A in HR 7775 The most striking di erence between the proðles of this line in s Lupi and HR 7775 is the clear presence of the A component of 199Hg in the latter star. If we adjust the overall mercury abundance to Ðt the damping wings of the line and WhiteÏs q-parameter to match this short-wavelength component, then we Ðnd a mercury abundance of log N(Hg) \ 6.44 (0.2 dex larger than in s Lupi), with q B 1.5 (Fig. 3). This gives about 80% of mercury in the form of 204Hg and about one part in a thousand as 199Hg. While the wings of the proðle Ðt well, there is again a relatively poor Ðt in the core of the line. Not only is the observed line core darker than the model proðle, but it is centered near the wavelength of 204Hg, rather than near the FIG. 1.È1942 A line in s Lupi. The solid line shows the observed GHRS spectrum (root name Z0G7010J). The dash-dotted line shows the combined rotationally broadened synthetic spectra for the primary and secondary stars, with appropriate instrumental broadening added as well. The dashed line shows the synthetic spectrum of the secondary alone, with the same instrumental broadening added. FIG. 3.È1942 A line in HR 7775 (root name Z365010A). Synthetic spectra are shown for log N(Hg) \ 6.44 with q \ 1.2 (dashed line, which shows extra absorption due to 199Hg) and q \ 1.8 (solid line).

8 No. 2, 1999 MERCURY IN s LUPI AND HR limitations to these assumptions could have a signiðcant e ect on the Ñux in the dark line cores. This suggests caution in attaching too much signiðcance to these discrepancies, especially when combined with the comparatively small number of observed counts and the correspondingly low S/N ratios in the line centers. FIG. 4.ÈDirect comparison between the observed GHRS spectra of the 1942 A Hg II line in s Lupi and HR The observed spectrum of s Lupi (dash-dotted line) has been shifted downward by the assumed continuum level of the secondary (as shown in Fig. 1), and scaled to the overall Ñux level of the observed spectrum of HR 7775 (solid line). In both stars, the deepest part of the line is centered at the wavelength of 204Hg. 202Hg line as predicted by the synthetic spectrum calculated with q \ 1.5. In Figure 4 we directly compare the observed data for the two stars. We can see that the darkest part of the line is centered at the same wavelength in both stars and is approximately the same depth if we assume the default secondary spectrum for s Lupi used in Figure 1. This argues against correcting the secondary Ñux in s Lupi as we did in Figure 2, but still allows for an equivalent change in the scattered-light correction. Removing the discrepancy between observed and calculated intensity in the depth of the line core of both stars would require reducing the scattered-light contribution from the default of 3.7% to about 2.5% (which is a plausible change given the uncertainties found by Cardelli et al. 1993), but the darkest part of the line in HR 7775 would still be centered at the wavelength of 204Hg. Fitting this would then require relative abundances of the isotopes that di er markedly from WhiteÏs formulation (e.g., 204Hg and 199Hg in the ratio given for q \ 1.5, but with other isotopes present with a fractional abundance of no more than a few times 10~3). This lack of an observable shift between the line core in these two stars also seems inconsistent with the large 202Hg fraction in HR 7775 found from the 3984 A line. It should be remembered that the line core is formed much higher in the atmosphere than the wings. At lower densities, any non-lte e ects will become larger. Even in the absence of signiðcant non-lte e ects, an atmospheric structure that is signiðcantly cooler in the outer layers than is our default model would naturally lead to a darker line core. Another possibility is vertical stratiðcation of the mercury abundance. A higher abundance in the coolest surface layers would also lead to a darker line core, and if in HR 7775 the extra mercury in such a surface layer were mostly 204Hg, then the center of the strongest lines might well need to be Ðtted with a di erent isotope mixture. The above discussion also assumes that the instrumental line broadening for SSA GHRS observations can be described by a simple Gaussian, and that there is no wavelength-dependent structure in the scattered light. Any A (6s 2S È6p 2P ) 1@2 3@2 Despite being the strongest of all mercury lines in s Lupi, the 1649 A resonance line does not provide any useful constraints on theoretical models beyond those provided by the 1942 A line. There is no part of the line proðle that is not substantially a ected by blending lines, and this renders impossible a straightforward Ðtting of the damping wings such as we did for the 1942 A line. The synthetic spectrum of this observation (Z2MB010M) also does a poorer job of reproducing the data than do the syntheses at longer wavelengths, and this makes appropriate continuum normalization more difficult. While we were fortunate in our observations of the 1942 A region in that rather featureless parts of the secondary spectrum were Doppler shifted to the wavelength of the Hg II line in the primary, this is not the case here. Even though the secondary contributes a smaller fraction of the continuum Ñux than at 1942 A, minor uncertainties in the secondaryïs absorption lines can easily obscure small proðle changes in the primaryïs spectrum. When combined with the greater line width and smaller isotope and hyperðne shifts in this transition, the above difficulties make it impossible to say much more than that the 1649 A Hg II line proðle in s Lupi is roughly consistent with the abundance and isotope mixture derived from the 1942 A line A (6p 2P È6d 2D ) 1@2 3@2 Since neither level in this transition involves 6s electrons, the isotope shift is quite small (only 2 ma between 204Hg and 198Hg, although the hyperðne structure of the 199Hg components is as large as 8 ma ). We anchor the wavelength scale in this region by using FTS wavelengths for Fe II lines at and A. Two high-s/n GHRS echelle spectra of s Lupi exist for this wavelength range. For the earlier of these two spectra, observation Z13J010M, proper modeling of the secondary features that are blended with the mercury line in the primary requires correcting the transition probabilities of the A Fe I and the A Fe II lines. Examination of these linesï contributions to the primary spectrum as seen in the GHRS data shows no trace of the A line and a strength for the A line that is much smaller than in the default synthetic spectrum. Once the transition probabilities of these lines have been corrected, the mercury line in the Z13J010M observation is essentially una ected by blending with secondary features, and we will use this observation for our primary analysis. In the other observation (Z13J510M) of this wavelength interval, the mercury line is blended with a complicated secondary feature. Our best-ðt proðle with q \ 3 requires an abundance of log N(Hg) \ 6.44, or 0.2 dex larger than that found for the 1942 A line. While there is clearly missing opacity between this Hg II line and a Si II line at A, the di erence in mercury abundance from our best Ðt to this line and from that to the 1942 A resonance line is scarcely larger than the expected relative uncertainties in the f-values.

9 950 PROFFITT ET AL. Vol. 512 FIG. 5.ÈSynthetic spectra near 1321 A are compared to the GHRS observation (Z2MB010B). As in Fig. 1, the solid line shows the observed data, and the dashed line shows the contribution of the secondary by itself. The dotted lines show the combined model Ñuxes for the primary and secondary for abundances in the primary of log n(hg) \ 6.24 and A [5d96s22D È5d96s6p(2D, 3P ) ] 5@2 3@2 2 5@2 This line is located in the red wing of the Tl II resonance line. Below 1500 A, few laboratory wavelength measurements accurate to 1 ma or better are available. The large number of unmodeled lines also makes setting the continuum level extremely difficult, but a reasonable normalization can produce an excellent Ðt to the blue wing of the two blended hyperðne components of the thallium line by using the abundance and isotopic mixture (pure 205Tl) found by Leckrone et al. (1996) from the 1908 A Tl II intercombination resonance line. Fitting this Hg II line then requires an abundance of log N(Hg) \ 5.69, which is 0.55 dex smaller than that determined from the 1942 A line (Fig. 5). The uncertainties in the continuum placement are unlikely to be large enough to explain this discrepancy; however, at this wavelength there are also substantial uncertainties in the continuum opacity that may signiðcantly a ect line formation. The upper energy level of this transition is a state with two excited electrons, and is at substantially higher energy than the upper states of the Hg II transitions we have discussed so far; it is possible that incompleteness of the model ion has a ected the transition probability calculation. In this regard, it is perhaps worth noting that the Cowan code calculations of both Sansonetti & Reader (1999) and Brage et al. (1999) Ðnd a substantially smaller transition probability (log gf \ and log gf \[0.03, respectively) than do the MCDF calculations of Brage et al. (1999) (log gf \ 0.419). A still smaller value, log gf B [0.13, would be needed to match the abundance derived from the 1942 A line. We suspect that the omission of the 5f levels from the MCDF calculations is adversely a ecting the solution for the entire set of 6s6p conðgurations. Unfortunately, a MCDF calculation that includes a set of conðgurations sufficiently large to address this question is, at present, computationally infeasible A [5d96s22D È5d96s6p(2D, 1P ) ] 5@2 5@2 1 7@2 This line is substantially more blended than the line at 1321 A, but shows similar discrepancies, requiring an abun- dance no larger than log N(Hg) \ 5.4 if the MCDF transition probability (log gf \ 0.351) is adopted. The upper level of this line is from the same conðguration as the 1321 A line, and its transition probability as calculated with a Cowan code program by Sansonetti & Reader (1999) (log gf \ 0.026) is also much smaller than the MCDF value. Given the good agreement between the 1942 and 1869 A lines, which involve only states with one excited electron, we consider difficulties with the transition probability calculations to be the most likely explanation for the problems encountered with this line and the one at 1321 A and 2253 A (6p 2P È6d 2D ) 3@2 3@2, 5@2 Two lines of this multiplet are covered by the low-s/n ACCUM observations made as part of a wavelengthcalibration program. For both lines, the MCDF calculations for the transition probabilities are within 0.1 dex of the Cowan code calculations. The line at A is less saturated than the A line and more sensitive to abundance variations. Both lines show good agreement with the abundance derived from the 1942 A line (see Table 7), while at least the A line is clearly inconsistent with the smaller abundances implied by the 1321 and 1331 A lines Hg I Previous work on optical lines of Hg I, especially the line near 4358 A (e.g., Wahlgren et al. 1994; Smith 1997) have found abundances similar to that found from the Hg II line at 3984 A. We have GHRS data for the ground-state lines at 1849 and 2536 A in s Lupi, but regrettably there is no GHRS data covering these lines in HR IUE data is of too low a resolution and too low a S/N ratio to give useful information on these lines. The intercombination line near 2536 A (6s21S È6s6p 3P ) is quite sensitive to abundance varia- 0 1 tions. Our wavelength alignment was made using two FTS measured lines: Ti II at A and Fe II at A. The observed line in s Lupi (Fig. 6) is clearly inconsistent with our nominal Hg II abundance, and requires an abundance D12 times smaller [log N(Hg) \ 5.1]. As is the case FIG. 6.ÈGHRS spectrum (Z0IX010G) near the Hg I 2536 A line. Synthetic spectra are shown for abundances of log n(hg) \ 5.0 and 6.2. Line styles are as in Fig. 5.

10 No. 2, 1999 MERCURY IN s LUPI AND HR for the 1942 A line, there is no evidence for any mercury isotope other than 204Hg. Much, though not all, of this ionization abundance anomaly would disappear if we omitted from the spectral synthesis of the secondary the absorption feature (Fe II A ) that is Doppler shifted so as to be coincident with the Hg I line in the primary. The relatively poor Ðt in the red wing of the 2537 A line might support such a change; however, the observed proðle of this line in the primary (which is visible in the same spectrum, although not shown in Fig. 6) suggests that the f-value of the iron line should actually be increased. Near 1849 A, we set the wavelength scale using the Fe II lines at and A and the Cr II line at A with FTS measured wavelengths. In Ðtting the 1849 A line, we Ðnd too much absorption from an Fe II line at A, even when no mercury line is included, and so we have reduced its gf-value from that of Kurucz by a factor of 10. This mercury line is rather saturated and relatively insensitive to abundance variations. The relatively low S/N ratio of this observation makes setting the continuum level difficult, but the line appears to be consistent with the abundance derived from the 2536 A line. Regardless of how the continuum level is set, Ðtting with an LTE abundance larger than log N(Hg) \ 5.6 is clearly excluded. The 6s21S È6s8p 1P transition at A is appar- 0 1 ently visible at about the expected strength in a GHRS spectrum (Z2MB0108). Unfortunately for this region of the spectrum, the large uncertainties in continuum levels and continuum opacities, and the many unidentiðed lines, together with the modest accuracy of available f-value determinations for this transition (Cowan code calculations only), prevent us from using this line to provide additional constraints on the nature of Hg I in s Lupi. Two other ground-state lines of 204Hg I at and A are also covered by existing GHRS data, but appear to be either too weak or too blended to be detectable, and, as is the case for the 1301 A line, available gf-values for these lines are of uncertain accuracy. An unidentiðed line near A in the GHRS spectra of s Lupi (observation Z2MB010I) and HR 7775 (observation Z32M0107) is probably not the mercury line, but may correspond to an unidentiðed weak line at A in the platinum atlas of Sansonetti et al. (1992). FIG. 7.ÈGHRS spectrum (Z0IX010R) near 1738 A in s Lupi, showing the Hg III lines. The synthetic spectra shown were calculated assuming no mercury (uppermost dotted line), and for log n(hg) \ 6.2, 6.7, and Line styles are as in Fig. 5. Hg II 1942 A (Fig. 7). The equivalent width in the feature near A is mostly due to an Fe III line, not the mercury line. We use a wavelength for this iron line ( A ) from Ekberg (1993), in place of the one from Kurucz. In this spectral region, we base our Ðnal wavelength calibration on two Fe II lines with FTS measured wavelengths of and A. The GHRS observation of the same region of the spectrum of HR 7775 is shown in Figure 8. The mercury ionization anomaly is much smaller, if not nonexistent, in this star (at most a factor of 2), and the central wavelength of the line is consistent with the same q-parameter that is needed to Ðt the 1942 A line. A direct comparison of this line in these two stars (Fig. 9) clearly shows the di erences in the line strength and the central wavelength Hg III All ground-state lines of Hg III have wavelengths shortward of the Lyman limit and are unobservable. The lines visible in the far-uv are all transitions from 5d96s levels with excitation energies between 5 and 8 ev to 5d96p levels with energies between 12 and 20 ev. Several of these lines are covered by our GHRS observations of s Lupi, but only a few appear to be sufficiently una ected by blending to be usefully studied A (5d96s 3D È5d96p 3P ) 2 2 The Hg III lines near A in s Lupi were Ðrst discussed by Leckrone et al. (1993), who found a large ionization anomaly relative to Hg II lines and also found that this line shows an isotope anomaly (essentially pure 204Hg) consistent with that found from Hg II lines. We Ðnd that the line near in the s Lupi spectrum is well Ðtted by q \ 3 and a mercury abundance of log N(Hg) \ 7.4, about 15 times larger than the nominal abundance determined from FIG. 8.ÈGHRS observation of HR 7775 (Z ) near 1738 A, compared to synthetic spectra computed with log N(Hg) \ 6.4 and 6.7 and an isotope mixture with q \ 1.2.