Mercury Elemental and Isotopic Abundances in Mercury-Manganese

THE ASTROPHYSICAL JOURNAL, 521 : 414È431, 1999 August 10 ( 1999. The American Astronomical Society. All rights reserved. Printed in U.S.A.

MERCURY ELEMENTAL AND ISOTOPIC ABUNDANCES IN MERCURY-MANGANESE STARS VINCENT M. WOOLF AND DAVID L. LAMBERT Department of Astronomy and McDonald Observatory, University of Texas at Austin, Austin, Texas 78712 ; vincent=astro.as.utexas.edu, dll=astro.as.utexas.edu Received 1999 January 21 ; accepted 1999 March 16

ABSTRACT Hg II abundances have been determined for 42 mercury-manganese (HgMn) stars by Ðtting synthetic spectra to observed spectra of the 3984 AŽ Hg II line. Twenty of the stars had lines sharp enough to allow their Hg isotopic abundance mixes to be estimated. The Hg abundance is reported for more HgMn stars here than in any other single work. No correlation was found between Hg II abundance and T or the eff mean central wavelength of HgMn j3984 stars. The mean central wavelength of j3984 , an indicator of the Hg isotopic mix, is loosely correlated with T : stars with primarily heavy Hg isotopes tend to be eff 204Hg and T in about the middle of the temcooler, although one star, 46 Aql, has almost pure perature range for HgMn stars. We Ðnd that there is no evidenceeff that any of the HgMn stars have 196Hg or 198Hg. For the very sharp-lined stars, the 204Hg abundance decreases with increasing T . No eff was correlation is seen between the mean central wavelength and the surface gravity. No correlation found between the projected rotational velocity and the Hg II abundance or the central wavelength of j3984, although this result may be biased by the selection of stars with low reported v sin i. Hg I j4358 was measured at high spectral resolution for seven HgMn stars. The isotopic shifts are too small, and the hyperÐne components are too weak to allow unambiguous isotopic abundance ratios to be found. Hg I abundances correlate fairly well with Hg II abundances. Some of the Hg isotopic mixtures are difficult to explain using only di†usion. HR 7245 has approximately equal abundances of 199Hg, 200Hg, 202Hg, and 204Hg but very little 201Hg, and 11 Per has Hg that is mostly 199Hg and 204Hg. Calculations show that hyperÐne splitting of 201Hg changes the radiative forces it feels compared with other isotopes, which may alter di†usion of that isotope enough to explain its absence in HR 7245, but we have found no possible explanation for the Hg isotopic mix found in 11 Per. These are the Ðrst very high resolution measurements of Hg II j3984 for HR 7245 and 11 Per. Although di†usion may be acting in HgMn stars, either there are one or more other mechanisms acting to help produce the overabundances and isotopic mixtures seen or our understanding of di†usion is lacking on some important point. Subject headings : di†usion È stars : abundances È stars : peculiar 1.

INTRODUCTION

& Landstreet 1993). W. P. Bidelman, an authority on the discovery of peculiar stars from spectroscopic scrutiny, declared ““ I do not now [see], nor have I ever seen, a normal A0 star ÏÏ (Eggen 1984). The common names for the CP stars reÑect the distinctive anomalies in the strengths of lines recognizable in optical spectra at low dispersion. For example, the mercurymanganese stars provide spectra with abnormally strong Mn II lines and a strong line of Hg II j3984. The unusual strength of Mn II lines was Ðrst noted by Morgan (1933). Bidelman (1962a, 1962b) later provided the correct identiÐcation of j3984. Hg abundances in HgMn stars have been found to be as large as about 400,000 times solar. Mn abundances have been found to be as large as about 600 times solar. In no case are the abundance anomalies restricted to the few elements for whom the class is named. Although the common names may reÑect a severe strengthening of lines, a marked weakening of lines of some elements may also occur ; for example, Co is generally underabundant by as much as a factor of 250 in HgMn stars exhibiting Mn overabundances of a factor of 100 or more (Smith & Dworetsky 1993a). Although other e†ects (e.g., departures from local thermodynamic equilibrium) may inÑuence the anomalous line strengths, it is widely supposed that the elementÏs abundance is the controlling inÑuence ; i.e., Mn and Hg are markedly overabundant in the HgMn stars. A detailed

The chemical composition of a main-sequence starÏs atmosphere is generally assumed to reÑect accurately the composition of its natal interstellar cloud. A few narrow exceptions to the assumption are generally recognized ; e.g., deuterium is predicted to be absent owing to its destruction, and the light elementsÈlithium, beryllium, and boronÈ may be depleted too. On the upper main sequence, striking broad exceptions to the assumption are common and are known as the chemically peculiar (CP) stars. Common names for di†erent classes of CP stars are the metallic-lined Am-Fm stars, the HgMn stars, the Bp-Ap stars (with j4200Si, Si, Si-Cr-Eu, Eu-Cr-Sr, and Sr types), the He-weak (with P-Ga, Sr-Ti, and Si types), the He-strong, and the j Bootis stars. Preston (1974) in a seminal review introduced a classiÐcation scheme : CP1 (Am-Fm), CP2 (Bp-Ap), CP3 (HgMn), and CP4 (He-weak). It must be stressed that the CP stars are not rare exceptions among upper main-sequence stars. For example, the late B-type main-sequence stars include the HgMn (CP3) with Wol† & Preston (1978) reporting that 16% of nonmagnetic late B-type stars and a remarkable 70% for slowly rotating (v ¹ 5 km s~1) stars are HgMn stars. The high frequency of CP stars is not restricted to the HgMn stars. Smith (1996) remarks that ““ nearly all slowly rotating A stars are chemically peculiar. ÏÏ Even apparently normal A stars show large star-to-star di†erences in composition (Hill 414

MERCURY IN HgMn STARS theoretical explanation for the abundance anomalies is lacking. In broad terms, the anomalies appear not to be due to alterations by nuclear processes but to a di†usive-like separation of elements. The site of the separation is not necessarily the stellar atmosphere, but the site and the atmosphere must be linked in order for the latter to be a†ected. The primary site of element separation could be exterior to the atmosphere, as in the proposals that invoke selective accretion by a magnetic star of elements from the interstellar medium (Havnes & Conti 1971 ; Havnes 1974) and that envisage accretion of gas but not dust from a dusty circumstellar shell (Venn & Lambert 1990). Alternatively, the site may be in the atmosphere, where gravitational settling competes with radiative acceleration in the absence of hydrodynamical mixing processes to create a di†usive separation between elements with a small cross section for photon absorption and elements with a large cross section. This competition between gravity and radiation may occur deeper in the star at the interface between the thin surface convection zone and the inner radiative envelope ; elements not supported there by radiation pressure di†use out of the convection zone, and other elements may be driven into the convection zone by radiation pressure. Di†usive separation, whether in or below the atmosphere, is reviewed by Michaud (1987, 1992) and Vauclair & Vauclair (1982). Progress toward a full understanding of the origins of chemical peculiarities will surely require additional observational evidence. In this paper, we present a survey of the isotopic abundances of mercury for HgMn stars. This was motivated in part by Vauclair & VauclairÏs (1982) suggestion that ““ Mercury manganese stars probably present the closest approach to the ideal situation of pure di†usion in stellar atmospheres. ÏÏ Several factors prompted this suggestion. First, very few HgMn stars have a detectable magnetic Ðeld in contrast to the Sr-Cr-Eu (CP2) stars where a Ðeld is commonly found. A magnetic Ðeld presumably complicates the operation of di†usive processes. Second, the stars are predicted to have a thin convective envelope. Third, single stars may become HgMn stars in sharp contrast to the case of the metallic line Am-Fm (CP1) stars, which are much more likely to be found in spectroscopic binaries than are ““ normal ÏÏ stars of a similar spectral class (Abt 1961, 1965). The fraction of HgMn stars found in binaries is about the same fraction found for normal stars (Schneider 1986). HgMn stars seem more likely to belong to spectroscopic binaries with periods of 4È16 days than do normal stars (Preston 1974), but HgMn stars are not found in binaries with periods shorter than 3 days (Smith 1996). HgMn stars in binaries do not appear di†erent from single HgMn stars, and, in some cases, both stars of a binary are HgMn stars. Fourth, many HgMn stars are slowly rotating. Rapid rotation may induce mixing that competes with other di†usive processes. The theory of rotationally induced mixing is not very well developed, so it makes sense to probe di†usive processes Ðrst in stars where such mixing is inoperative. Rapid rotation can have indirect e†ects, too, on the di†usive processes. Michaud (1982) calculated that meridional circulation in stars rotating in excess of 90 km s~1 may sustain the He II convection zone in stars with the e†ective temperature of the HgMn stars. This convection zone below the outer thin convection zone and the radiative envelope is predicted to disappear as He di†uses out the bottom of the zone in more slowly rotating stars, and its disappearance

415

greatly inÑuences the predicted composition of the atmosphere. In this paper, we present measurements of the isotopic abundances of Hg from a high-resolution spectroscopic survey of northern HgMn stars and review the results in light of predictions from models invoking the di†usive separation of elements. Although determinations of the elemental Hg abundance have been reported previously for many HgMn stars from optical (e.g., Cowley & Aikman 1975 ; White et al. 1976 ; Heacox 1979 ; Smith 1997 ; Adelman, Ryabchikova, & Davydova 1998 and references therein) and ultraviolet (Leckrone 1984 ; Smith 1997) spectra, the only previous works we know of with a spectral resolution adequate to resolve the isotopic structure of optical Hg II lines are the pioneering study by White et al. (1976) and a recent preliminary report for Ðve stars by Hubrig, Castelli, & Mathys (1998). Proffitt et al. (1999) observed several ultraviolet Hg I, Hg II, and Hg III lines in the HgMn stars s Lup and HR 7775. They report that for s Lup, all observed lines are consistent with a pure 204Hg mixture and that HR 7775 shows evidence of some lighter isotopes. Our paper reports Hg abundances for 42 HgMn stars and Hg isotopic mixtures for 20 stars, more than any previous study. The spectral resolution used for the isotopic mixture determination (j/*j [ 160,000) is greater than reported in any previous work on HgMn stars. The range of Hg abundances in HgMn stars determined in our study is similar to that found by others. The Hg abundances can range from as large as about 400,000 times the solar abundance to small enough that Hg II j3984 is undetectable. 2.

OBSERVATIONS AND DATA REDUCTION

Stars selected for observation were known HgMn stars north of decl. \ [30¡ with reported v sin i \ 35 km s~1. Two Mn-Hg star catalogs (Schneider 1981 ; Renson, Gerbaldi, & Catalano 1991) were especially helpful in the star selection. Observations of Hg II j3984 were made at McDonald Observatory using the 2.7 and 2.1 m telescopes. Spectra were measured on the 2.7 m telescope using the crossdispersed echelle spectrometer at the coude focus (Tull et al. 1995). Resolution depended on at which focus of the spectrometer the detector was placed. One set of measurements was made at the focus that gives a resolution of 60,000 and provided complete wavelength coverage between 3600 and 6900 AŽ . The others were made at very high resolution focus. For all observations at the very high resolution focus, resolution fell in the range 160,000 [ j/*j [ 190,000. The spectra taken at the very high resolution focus include 12 spectral orders in the wavelength range from about 3570 to 4030 AŽ . Each order covers about 15 AŽ . There are approximately 30 AŽ gaps in coverage between the orders. Spectra were measured on the 2.1 m telescope using the Sandiford Cassegrain echelle spectrograph (McCarthy et al. 1993). The 2.1 m spectra covered a wavelength range from about 3950 to 4300 AŽ with no gaps in coverage. Resolution was approximately 60,000. For all observations the detector was a large-format CCD. Signalto-noise ratio values for the Hg II spectra were generally between 100 and 300. Observations of Hg I j4358 were made using the coude spectrograph of the 2.7 m telescope at very high resolution focus (j/*j B 180,000). Signal-to-noise ratio values for the Hg I spectra were between 400 and 500.

416

WOOLF & LAMBERT

For each nightÏs data overscan, bias, and scattered light subtraction, Ñat lamp combination and normalization, division by Ñat lamp, spectral reduction, and thorium-argon spectrum wavelength calibration were performed using standard IRAF packages. Wavelength shifts for radial velocity corrections were determined by measuring the wavelengths of a large number of atomic lines and taking the average of the velocity shift needed to place them at their rest wavelengths. If multiple spectra of the same star were made in one night, they were co-added. For one star, the short-period double-lined binary 74 Aqr (P \ 3.43 days), the velocity correction had to be applied to correct for binary orbital motion before co-adding spectra. 3.

ANALYSIS

3.1. Stellar Atmospheric Parameters For most program stars T and log g were determined eff using StroŽmgren uvby b photometry1 (Mermilliod, Mermilliod, & Hauck 1997) and a FORTRAN program (Moon 1985 ; Napiwotzki, SchoŽnberner, & Wenske 1993) that interpolates the grids of Moon & Dworetsky (1985). For stars where more reliable stellar parameters were available, the better data were used. In particular, parameters for 66 Eri A and B were taken from Yushchenko et al. (1998), HR 4072 from Smith (1997), • CrB from Adelman (1989), 46 Dra A and B from Adelman et al. (1998), and 112 Her from Ryabchikova, Zakharova, & Adelman (1996). Except for • CrB and 46 Dra B, the parameters calculated from photometry agree within the errors with the parameters that were used : using photometry is not the best method for determining T and log g in unresolved binaries when light from eff the secondary star cannot be neglected. Parameters for program stars are listed in Table 1. All stars listed are HgMn stars except 66 Eri A, which was included for comparison with the HgMn star 66 Eri B. Solar abundance LTE model atmospheres with the speciÐed stellar parameters were calculated using the ATLAS9 model atmosphere program (Kurucz 1993). Several of the program stars are in spectroscopic binaries. Corrections were made for continuum Ñux due to the companion star by subtracting the estimated fraction of the Ñux from the spectrum and renormalizing before further analysis. Table 2 lists the binary corrections made for programs stars and the sources of companion Ñux estimates. Corrections made using Bright Star Catalog (Hoffleit & Warren 1991) and Hipparcos (ESA 1997) data are based on magnitude di†erences between components. Microturbulences were determined by calculating Fe II and in some cases Cr II abundances from equivalent widths using the stellar abundance program MOOG (Sneden 1973). For our study, MOOGÏs treatment of hydrogen lines was improved, and additional line-broadening sources such as Van der Waals broadening were added. Microturbulence values were varied until there was no slope in Fe and/or Cr abundances versus equivalent width. For Ðve stars, microturbulence values from other works were used, as indicated in Table 1. Rotational velocities were determined for stars measured at very high resolution and for stars with high enough rotational broadening that v sin i could be easily measured with the lower resolution spectra. For these stars v sin i was 1 The photometry database of Mermilliod et al. (1997) can be searched at the following Web site : http ://obswww.unige.ch/gcpd/gcpd.html.

Vol. 521

determined by varying abundances and v sin i for synthetic spectra of unblended Fe and Si lines using MOOG until there was a good Ðt to observed spectra. The quoted uncertainties in v sin i values reÑect the range for which the Ðt of synthetic to observed spectra was reasonable. Values of v sin i thus determined are reported in Table 5 below. 3.2. Hg II Abundance Determinations Hg II abundances were estimated by Ðtting synthetic spectra calculated by the LTE spectral synthesis program MOOG to observed stellar spectra. The FORTRAN program AMOEBA from Press et al. (1992), which uses the downhill simplex method to Ðnd minima or maxima of a function, was added to MOOG to allow automatic Ðtting of synthetic spectra to observed spectra by minimizing the s2 di†erence between them. The synthetic spectra were Ðtted to a region 0.8È2.1 AŽ wide around Hg II j3984, depending on the starÏs rotational broadening. Hg II isotopic and hyperÐne component wavelengths were determined using Fourier transform spectrometer (FTS) measurements by the atomic spectroscopy group at Lund University (U. Litzen 1997, private communication) and data from Proffitt et al. (1999). The gf-value for the Hg II line was taken from Proffitt et al. (1999). The line list used is shown in Table 3. The contributions of Cr I j3983.896 and Fe I j3983.956 were included in the synthetic spectra. The Cr and Fe abundances used in the synthetic spectra were calculated from equivalent widths of Cr and Fe lines that fell in the spectral orders around the Hg II line and are indicated in Table 5. When Ðtting a synthetic spectrum to an observed spectrum, the continuum level was set by hand, the Cr and Fe abundances were held constant, and v sin i and the Hg isotopic abundances were allowed to vary until the s2 di†erence between the spectra was minimized. Fitting was not perfect. Several runs were made using di†erent starting Hg abundances, and the results of the best Ðts were averaged. For sharp-lined stars like HR 7245 where the isotopic components are resolved in the very high resolution spectra, isotopic abundance ratios could be determined unambiguously. For some stars with lines that are not quite so sharp, isotopic abundance fractions could be estimated, but with higher uncertainties. Uncertainties were estimated by running the automatic synthetic spectrum Ðtting program several times and Ðnding the standard deviation of the isotopic fraction results. For stars with a higher v sin i, a range of isotopic abundances could give equally good Ðts, but the central wavelength of the 3984 AŽ feature is still a good indicator of how much the Hg II abundance is shifted toward heavy or light isotopes. The central wavelengths are included in Table 5. Central wavelengths of the line were found using the SPLOT package of IRAF. The reported values represent the mathematical, integrated centroids of the line. These may di†er somewhat from the central wavelengths given by White et al. (1976), which, as noted in Smith (1997), were determined by eye from photographic plates by centering the crosshair of a traveling microscope on the 3984 AŽ feature. Stars with very weak 3984 AŽ features do not have central wavelengths or equivalent widths listed in the table. Fourteen stars observed for this paper were included in Smith (1997). Figure 1 compares the Hg II abundances2 and 2 I use the notation A(X) 4 log

10

(N /N ) ] 12.00. X H

No. 1, 1999

MERCURY IN HgMn STARS

417

TABLE 1 PARAMETERS FOR PROGRAM STARS

HD

HR

Name

T eff (K)

log g (cgs)

m used (km s~1)

m Source a

3.54 4.19 4.19 4.25 3.77 3.77 3.77 3.40 3.77 3.80 3.36 3.28 3.87 3.72 3.63 3.81 3.80 4.02 3.31 3.75 3.97 3.65 3.73 4.19 3.97 3.83 3.97 4.01 3.98 4.01 4.00 4.10 4.10 3.44 3.56 3.74 3.43 4.12 3.34 3.56 4.05 4.05

0.6 0.8 1.0 0.9 1.6 0.5 0.0 2.0 0.5 0.5 1.1 1.0 1.0 2.0 0.1 0.0 0.2 1.5 1.2 1.8 0.0 0.9 1.0 1.0 0.4 0.5 1.3 0.5 2.0 0.0 1.9 1.9 0.5 1.0 0.7 0.6 0.5 0.0 1.0 2.0 0.0 1.0

1 1 1 1 1 2 1 1 1 2 1 1 1 1 2 1 2 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 3 1 1 1 1 1 1 1 1 1

4.25

0.9

1

HgMn Stars 4382 . . . . . . . . . . . 16727 . . . . . . . . . . 27295 . . . . . . . . . . 32964 B . . . . . . . 33904 . . . . . . . . . . 35548 . . . . . . . . . . 49606 . . . . . . . . . . 53244 . . . . . . . . . . 53929 . . . . . . . . . . 58661 . . . . . . . . . . 63975 . . . . . . . . . . 70235 . . . . . . . . . . 72208 . . . . . . . . . . 75333 . . . . . . . . . . 77350 . . . . . . . . . . 78316 . . . . . . . . . . 89822 . . . . . . . . . . 101391 . . . . . . . . 106625 . . . . . . . . 110073 . . . . . . . . 129174 . . . . . . . . 143807 . . . . . . . . 144206 . . . . . . . . 144661 . . . . . . . . 145389 . . . . . . . . 149121 . . . . . . . . 159082 . . . . . . . . 169027 . . . . . . . . 172728 . . . . . . . . 172883 . . . . . . . . 173524 A . . . . . . 173524 B . . . . . . 174933 . . . . . . . . 178065 . . . . . . . . 182308 . . . . . . . . 186122 . . . . . . . . 190229 . . . . . . . . 193452 . . . . . . . . 202671 . . . . . . . . 207857 . . . . . . . . 216494 . . . . . . . . 220933 . . . . . . . .

208 785 1339 1657 B 1702 1800 2519 2657 2676 2844 3059 3273 3361 3500 3595 3623 4072 4493 4662 4817 5475 5971 5982 5998 6023 6158 6532 ... 7018 7028 7049 A 7049 B 7113 7245 7361 7493 7664 7775 8137 8349 8704 8915

23 Cas 11 Per 53 Tau 66 Eri B k Lep ... 33 Gem c CMa ... ... f CMi ... ... 14 Hya l Cnc i Cnc ... ... c Crv l Cen n1 Boo • CrB t Her ... / Her 28 Her ... 38 Dra ... ... 46 Dra A 46 Dra B 112 Her ... ... 46 Aql ... ... 30 Cap ... 74 Aqr 69 Peg

13,200 14,550 12,000 10,900 12,750 11,000 14,250 13,600 13,900 13,450 13,500 12,350 10,900 12,250 10,300 13,650 10,650 12,500 12,000 12,900 13,050 11,250 12,000 15,800 11,800 10,900 11,100 11,500 10,700 11,300 11,700 11,100 13,100 12,350 13,550 12,900 13,250 10,700 13,700 13,300 11,950 10,950

66 Eri BÏs Companion 32964 A . . . . . . .

1657 A

66 Eri A

11,100

a Microturbulence sources : (1) This paper ; (2) Smith 1997 ; (3) Ryabchikova et al. 1996.

equivalent widths of the 3984 AŽ line from this paper and SmithÏs. The equivalent widths from this paper tend to be slightly larger, with one of the largest di†erences [data point on lower panel at (135, 85)] being caused by the binary Ñux correction applied for i Cnc in this work. For the other large di†erence [data point at (177, 144)] we simply measured a larger equivalent width than that found in White et al. (1976) for HR 2844. The di†erences in A(Hg) come from di†erences in spectra, binary corrections, analysis programs, and stellar parameters used. For several stars, the larger microturbulence used in this paper produced smaller calculated Hg II abundances than found in Smith (1997). Figure 2 compares the central wavelengths of the 3984 AŽ

lines in this work and Smith (1997). The agreement is fairly good. 3.3. Hg I Abundance Determinations Hg I j4358, the strongest Hg I line in HgMn stars, was observed in seven stars using the 2.7 m telescopeÏs coude spectrograph at very high resolution to see if using a higher resolution than used in other studies would allow isotopic abundance determinations for Hg I. Di†erences in the Hg abundances and isotopic mixtures measured using Hg I j4358 and Hg II j3984 are predicted by some models in which Hg isotopes are concentrated in thin layers in the atmosphere. Previous calculations of Hg abundances from

418

WOOLF & LAMBERT

Star 23 Cas . . . . . . . . . 11 Per . . . . . . . . . 53 Tau . . . . . . . . 66 Eri A . . . . . . . 66 Eri B . . . . . . . k Lep . . . . . . . . . . HR 1800 . . . . . . 33 Gem . . . . . . . c CMa . . . . . . . . . HR 2676 . . . . . . HR 2844 . . . . . . f CMi . . . . . . . . . HR 3273 . . . . . . HR 3361 . . . . . . 14 Hya . . . . . . . . l Cnc . . . . . . . . . . i Cnc . . . . . . . . . . HR 4072 . . . . . . HR 4493 . . . . . . c Crv . . . . . . . . . . l Cen . . . . . . . . . . n1 Boo . . . . . . . . • CrB . . . . . . . . . . t Her . . . . . . . . . . HR 5998 . . . . . . / Her . . . . . . . . . 28 Her . . . . . . . . . HR 6532 . . . . . . 38 Dra . . . . . . . . HR 7018 . . . . . . HR 7028 . . . . . . 46 Dra A . . . . . . 46 Dra B . . . . . . 112 Her . . . . . . . HR 7245 . . . . . . HR 7361 . . . . . . 46 Aql . . . . . . . . . HR 7664 . . . . . . HR 7775 . . . . . . 30 Cap . . . . . . . . HR 8349 . . . . . . 74 Aqr . . . . . . . . . 69 Peg . . . . . . . . .

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TABLE 2

TABLE 3

BINARY TREATMENT FOR PROGRAM STARS

LINES FOR Hg II SYNTHETIC SPECTRA

Binary Type SB Single SB SB2 SB2 ... Double ... ... ... Double Single Single SB2 ... SB1 SB1 SB2 Single SB SB1 Double SB2 Single Single SB1 Single SB1 Single Single Single SB2 SB2 SB2 SB1 Double Single SB1 Double Double ? SB1 SB2 ...

Correction ([) None None None 0.487 0.513 none 0.20 None None None None None None 0.075 None None 0.35 0.143 None None None None 0.20 None None None None None None None None 0.37 0.63 0.102 None None None None 0.05 None None 0.32 None

Correction Sourcea ... ... ... Yus Yus ... Hipparcos ... ... ... Hipparcos ... ... Mass ratio ... ... BSC Smith ... ... ... Hipparcos Smith ... ... ... ... ... ... ... ... ARD ARD RZA ... ... ... ... Hipparcos ... ... Mass ratio ...

a Mass ratio uses log (L /L ) \ 3.99 log (M /M ) : BSC \ Hoffleit & 1 ESA 2 1997 ; Smith1 \ 2Smith 1997 ; RZA \ Warren 1991 ; Hipparcos \ Ryabchikova et al. 1996 ; ARD \ Adelman et al. 1998 ; Yus \ Yushchenko et al. 1997.

j4358 (Smith 1997 ; Wahlgren, Adelman, & Robinson 1994) in HgMn stars have treated the line as resulting from a single isotope, which gives the Hg abundance when the weak line is present in the cooler HgMn stars but does not give any isotopic information. The isotopic splitting of j4358 is not as large as the splitting of j3984. While the hyperÐne components of the odd isotopes are spread out enough to be detected at this resolution, they were not strong enough in any of the observed stars to be detected unambiguously, even using spectra with signal-to-noise ratios of 400È500. Radial velocity corrections for the stars were made by measuring the wavelengths of as many unblended lines as possible for each star and applying the average velocity shift indicated by these lines. Unfortunately, this method seemed to have a

Species 196Hg II . . . . . . 198Hg II . . . . . . 199Hg II . . . . . . Cr I . . . . . . . . . . 200Hg II . . . . . . 201Hg II . . . . . .

Fe I . . . . . . . . . . 202Hg II . . . . . . 204Hg II . . . . . .

j (AŽ )

s (eV)

log gf

3983.778 3983.840 3983.834 3983.841 3983.855 3983.896 3983.912 3983.927 3983.936 3983.939 3983.940 3983.941 3983.946 3983.947 3983.956 3983.993 3984.072

4.40 4.40 4.40 4.40 4.40 2.54 4.40 4.40 4.40 4.40 4.40 4.40 4.40 4.40 2.73 4.40 4.40

[1.529 [1.529 [2.906 [1.955 [1.763 0.300 [1.529 [1.955 [2.729 [2.412 [2.162 [2.468 [2.643 [3.751 [1.001 [1.529 [1.529

random error of about ^0.005 AŽ , as shown by the star-tostar di†erences in the position of Y II j4358.9 after velocity correction. This uncertainty is about the size of the isotope shifts between consecutive even Hg I isotopic components for j4358. Therefore, total Hg I abundances could be calculated using the same model atmospheres and synthetic synthesis program as was used to Ðnd Hg II abundances, but isotopic ratios could not. However, if the Y II line is used to anchor the wavelength scale, the central wavelengths of the Hg I lines can be used to compare the relative isotopic mixtures : a shorter central wavelength means heavier isotopes. The line list used is shown in Table 4. Hg I atomic data were taken from Smith (1997), and other atomic data were taken from the atomic line list of R. L. Kurucz. The same method used to Ðnd Hg II abundances was used to Ðnd the Hg I abundances : v sin i and Hg I abundances were allowed TABLE 4 LINES FOR Hg I SYNTHETIC SPECTRA

Species Fe II . . . . . . . . . 199Hg I . . . . . . 201Hg I . . . . . .

199Hg 201Hg 204Hg 202Hg 200Hg 198Hg 196Hg 201Hg

I...... I...... I...... I...... I...... I...... I...... I......

199Hg I . . . . . . 201Hg I . . . . . . Fe I . . . . . . . . . . 199Hg I . . . . . . Y II . . . . . . . . . .

j (AŽ )

s (eV)

log gf

4358.162 4358.175 4358.227 4358.240 4358.288 4358.315 4358.316 4358.320 4358.326 4358.332 4358.337 4358.341 4358.354 4358.364 4358.379 4358.444 4358.501 4358.519 4358.728

9.44 4.89 4.89 4.89 4.89 4.89 4.89 4.89 4.89 4.89 4.89 4.89 4.89 4.89 4.89 4.89 2.95 4.89 0.10

[1.686 [1.424 [1.294 [1.324 [2.018 [0.725 [1.821 [0.469 [0.469 [0.469 [0.469 [0.469 [0.924 [1.324 [1.115 [1.294 [2.247 [1.425 [1.320

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MERCURY IN HgMn STARS

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7 4.04

4.02

A(Hg) (Smith)

6.5

4

6 3.98

3.96

5.5

3.94

5 5

5.5

6 A(Hg) (this paper)

6.5

7

3.92 3.92

3.94

3.96

3.98

4

4.02

4.04

FIG. 2.ÈComparison of central wavelength of j3984 from this paper and Smith (1997)

150

100

50

50

100

150

FIG. 1.ÈHg II abundances and 3984 AŽ line equivalent widths from this paper and Smith (1997) compared.

to vary in the auto-Ðtting routine to Ðnd the best Ðt to the observed spectrum. This method did not yield usable Hg isotopic abundances. For example, synthetic spectra of Hg I j4358 of • CrB can be created that Ðt the observed spectrum well with a 202Hg fraction anywhere between 0% and 75% by altering the 200Hg and 204Hg abundances to compensate. 4.

RESULTS

4.1. Hg II Abundances Spectra of program stars are shown in Figures 3aÈ3d, arranged in the order of increasing 3984 AŽ equivalent widths. Note that the scale for the Ñux changes from panel to panel. Hg II abundances for program stars are listed in Table 5. Note that the log gf-value used for the Hg abundance determination ([1.529 ; Proffitt et al. 1999) is larger than that

used in most other abundance studies ([1.730). So absent any other di†erences, the abundances reported here will be 0.20 dex smaller. The solar system mercury abundance as determined from meteorites is 1.17 (Grevesse, Noels, & Sauval 1996). Hg abundances reported here range from undetectable (less than 2.0 dex larger than solar) to 5.6 dex larger than the solar value. Hg abundances are reported here for more HgMn stars (42) than in any other single work. In Table 5 it is also noted for each star whether the 3984 AŽ region was observed at very high resolution (j/*j B 160,000). A few of the stars that were observed only at the lower resolution (j/*j B 60,000) may reveal sharper lines if observed at higher resolution. There are several sources of possible error in the abundance determinations, only some of which are included in the stated uncertainties in Table 5. For the program stars, uncertainties in photometry normally give uncertainties in T in the range of about [100 to 200 K and uncertainties inefflog g of about ^0.03 dex. A 200 K error in T will give errors in the Hg II abundance of about 0.02 dex.eff A 0.03 dex error in log g will give errors in the Hg II abundance of about 0.01 dex. Uncertainties in the determined Cr and Fe abundances used in synthetic spectra give an uncertainty in the Hg II abundance of about 0.01 dex. Uncertainties in microturbulence vary quite a bit. For sharp-lined stars, where it was simple to determine if the Fe and Cr lines used to Ðnd microturbulence were blended and it was easier to determine the continuum level, microturbulence uncertainties cause Hg II abundance uncertainties of about 0.03 dex. These errors sum to give a Hg II abundance uncertainty of about 0.07 dex. For stars with broader lines, microturbulence determination is less precise and can lead to a Hg II abundance uncertainty from j3984 about 0.35 dex if the line is very strong. In calculating the Hg II abundance uncertainty reported, the 0.07 dex value was used as the part due to uncertainties in stellar parameter for all stars. There are additional uncertainties in A(Hg) due to uncertainties in Ðtting synthetic spectra to observed spectra,

420

WOOLF & LAMBERT

Vol. 521

53 Tau 3

HR 1846

11 Per

30 Cap

HR 3361

2 HR 1883

2.5

74 Aqr

HR 2676

2

HR 7664

23 Cas 1.5

l Cen

46 Aql 1.5

28 Her 1

HR 7028

1

33 Gem

HR 7245

a

3983

b 3983.5

3984

3984.5

3983

46 Dra B 4

4 HR 5998 66 Eri B

3983.5

3984

3984.5

3984

3984.5

69 Peg HR 4072 HR 7361

HR 6532 3

HR 7775

3 14 Hya

HR 1800

HR 7018

HR 3273 2

46 Dra A

2 38 Dra HR 8349

112 Her 1

1

HR 4493

HR 2844

c 3983

d 3983.5

3984

3984.5

3983

3983.5

FIG. 3.ÈObserved stellar spectra around Hg II j3984, corrected for radial velocities and binary companion Ñux contributions. Flux (vertical) scale is di†erent in separate panels.

including noise in the observed spectrum, possible errors in the placement of the continuum, and wavelength corrections due to radial velocities. These uncertainties were estimated by Ðtting synthetic spectra for each star several times and varying the continuum placement and wavelength shift within the values reasonable for each spectrum. For broader lined stars, some of the abundance uncertainty came from di†erent possible isotope ratios that Ðtted the

observed spectrum equally well. The uncertainties estimated by this method were added to the 0.07 dex error estimated from uncertainties in stellar parameters to give the uncertainty values listed. Not included in the uncertainty estimates were the e†ects of the up to 0.1 dex possible systematic underestimate of log g due to the He deÐciency found in many HgMn stars (Smith & Dworetsky 1993a) or any uncertainty in stellar

No. 1, 1999

MERCURY IN HgMn STARS

421

TABLE 5 Hg II ABUNDANCES FOR PROGRAM STARS

Star

A(Cr)

used

A(Fe)

used

A(Hg II)

Hg Central j (3980 AŽ ] )

W j (mAŽ )

v sin i (km s~1)

R B 160,000 ?

17 29 ... 78 116 96 56 19 6.5 177 10 89 35 119 44 135 106 96 28 54 150 90 111 77 66 28 78 134 125 51 90 71 93 66 83 17 53 79 ... 152 43 98

23.0 ^ 2.0 4.8 ^ 0.4 6.0 ^ 2.0 ... ... 2.5 ^ 0.5 ... ... ... 27.0 ^ 2.0 28.0 ^ 3.0 20.0 ^ 3.0 ... 35.0 ^ 3.0 ... 7.3 ^ 0.5 3.2 ^ 0.3 54.0 ^ 3.0 ... ... 14.0 ^ 2.0 \1.3 10.0 ^ 0.5 47.0 ^ 4.0 9.0 ^ 1.0 9.0 ^ 0.5 ... 32.0 ^ 4.0 ... ... ... ... ... \1.0 8.0 ^ 0.5 \1.0 9.0 ^ 0.4 1.8 ^ 0.3 ... 12.8 ^ 0.5 3.5 ^ 0.5 35. ^ 2.0

y y y y n y n n n n y y n n n y y n n n y y y n y y n n n n n n n y y y y y n y y n

...

...

y

HgMn Stars 23 Cas . . . . . . . . . 11 Per . . . . . . . . . 53 Tau . . . . . . . . 66 Eri B . . . . . . . k Lep . . . . . . . . . . HR 1800 . . . . . . 33 Gem . . . . . . . c CMa . . . . . . . . . HR 2676 . . . . . . HR 2844 . . . . . . f CMi . . . . . . . . . HR 3273 . . . . . . HR 3361 . . . . . . 14 Hya . . . . . . . . l Cnc . . . . . . . . . . i Cnc . . . . . . . . . . HR 4072 . . . . . . HR 4493 . . . . . . c Crv . . . . . . . . . . l Cen . . . . . . . . . . n1 Boo . . . . . . . . • CrB . . . . . . . . . . t Her . . . . . . . . . . HR 5998 . . . . . . / Her . . . . . . . . . 28 Her . . . . . . . . . HR 6532 . . . . . . 38 Dra . . . . . . . . HR 7018 . . . . . . HR 7028 . . . . . . 46 Dra A . . . . . . 46 Dra B . . . . . . 112 Her . . . . . . . HR 7245 . . . . . . HR 7361 . . . . . . 46 Aql . . . . . . . . . HR 7664 . . . . . . HR 7775 . . . . . . 30 Cap . . . . . . . . HR 8349 . . . . . . 74 Aqr . . . . . . . . . 69 Peg . . . . . . . . .

6.03 5.30 5.67 6.44 5.56 5.39 6.10 5.82 6.18 6.30 6.00 5.67 5.67 5.56 6.50 6.11 6.38 5.86 5.07 6.63 6.64 6.02 6.10 6.40 6.70 5.90 5.67 6.23 5.79 5.88 5.70 6.09 5.48 5.82 5.30 5.10 5.10 6.30 5.67 5.89 5.82 5.67

7.86 7.59 7.50 7.64 7.39 7.27 8.00 7.65 8.01 7.20 7.91 7.50 7.50 7.39 7.65 7.99 7.87 7.67 6.94 7.00 7.29 7.85 7.20 8.25 7.65 7.60 7.50 8.06 7.70 7.71 7.69 7.72 8.43 7.32 7.85 7.88 8.10 7.95 7.88 7.35 7.65 8.10

4.41 ^ 0.30 5.06 ^ 0.14 [3.55 5.83 ^ 0.16 5.62 ^ 0.17 6.36 ^ 0.21 5.32 ^ 0.14 4.50 ^ 0.17 4.11 ^ 0.14 6.75 ^ 0.33 4.18 ^ 0.22 5.58 ^ 0.32 5.11 ^ 0.09 5.67 ^ 0.11 5.48 ^ 0.28 6.56 ^ 0.31 6.64 ^ 0.09 6.11 ^ 0.33 4.58 ^ 0.11 5.04 ^ 0.10 6.72 ^ 0.25 5.74 ^ 0.17 5.82 ^ 0.13 6.08 ^ 0.39 5.81 ^ 0.14 5.31 ^ 0.22 5.59 ^ 0.11 6.42 ^ 0.28 6.04 ^ 0.08 5.61 ^ 0.12 5.75 ^ 0.12 5.90 ^ 0.16 6.14 ^ 0.09 5.13 ^ 0.09 5.42 ^ 0.14 4.77 ^ 0.14 5.18 ^ 0.28 6.43 ^ 0.11 [4.00 5.88 ^ 0.10 5.20 ^ 0.10 6.26 ^ 0.11

3.944 3.991 ... 3.974 3.938 3.999 3.944 3.983 3.905 3.940 3.941 3.985 4.026 3.970 4.043 3.941 4.012 3.952 3.991 3.946 3.948 3.978 3.974 4.003 4.006 4.068 3.954 3.945 3.988 3.994 4.013 4.037 3.984 3.948 3.936 4.073 3.963 4.021 ... 3.960 3.933 4.011

66 Eri BÏs Companion 66 Eri A . . . . . . .

5.90

7.70

[2.2

parameters or model atmospheres due to the e†ects of the HgMn starsÏ unusual chemical compositions on metal opacities in the atmospheres. Also not included are Hg II abundance uncertainties due to possible errors in binary Ñux corrections. Estimates of Hg II isotopic fractions in sharper lined HgMn stars are listed in Table 6 and compared with the terrestrial isotopic mixture. Stars where the Hg isotopic components of j3984 were not resolved have larger uncertainties in their estimated isotopic mixtures. Figure 4 shows the observed spectra of the very sharp lined stars with Hg isotopic and hyperÐne components indicated. Uncertainties in isotopic ratios were estimated by making several synthetic spectrum Ðts and varying continuum and wavelength shifts within values reasonable for each sharp-lined star.

...

4.2. Hg I Abundances Line j4358 is the strongest Hg I line in the visible in HgMn stars, and it appears that it is free of blending. Still, it is rather weak, even in cooler HgMn stars where there is more Hg I. Hg I abundances are listed in Table 7 for the seven stars for which the 4358 AŽ region was observed at very high resolution. The uncertainties were found by adding the uncertainties due to the spectrum Ðtting routine to those due to stellar parameter uncertainties (^0.01 dex from m, ^0.02 dex from log g, and ^0.06 dex from T ). eff The wavelength calibration was not precise enough to allow Hg isotopic composition estimates to be made using the small isotope splitting of j4358. The wavelength of the line does indicate somewhat whether the Hg I abundance is

422

WOOLF & LAMBERT

Vol. 521

TABLE 6 Hg ISOTOPE FRACTIONS FOR SHARPER LINED STARS Star

196Hg

198Hg

199Hg

200Hg

201Hg

202Hg

204Hg

Sun . . . . . . . . . . . .

0.0015

0.100

0.169

0.231

0.132

0.298

0.0685

0.008 ^ 0.013 0.076 ^ 0.009 0.006 ^ 0.001 0.078 ^ 0.003 0.004 ^ 0.004 0.042 ^ 0.024 0.002 ^ 0.003 0.005 ^ 0.001 0.155 ^ 0.012

0.012 ^ 0.019 0.533 ^ 0.033 0.376 ^ 0.052 0.505 ^ 0.006 0.027 ^ 0.018 0.239 ^ 0.012 0.012 ^ 0.002 0.367 ^ 0.002 0.236 ^ 0.007

0.641 ^ 0.038 0.373 ^ 0.040 0.611 ^ 0.052 0.354 ^ 0.004 0.965 ^ 0.018 0.209 ^ 0.018 0.961 ^ 0.010 0.624 ^ 0.002 0.109 ^ 0.012

0.268 ^ 0.056 0.255 ^ 0.106 0.487 ^ 0.138 0.326 ^ 0.104 0.340 ^ 0.164 0.379 ^ 0.194 0.138 ^ 0.178 0.505 ^ 0.147 0.122 ^ 0.054 0.616 ^ 0.214 0.315 ^ 0.197

0.067 ^ 0.046 0.302 ^ 0.199 0.239 ^ 0.085 0.558 ^ 0.135 0.164 ^ 0.056 0.533 ^ 0.192 0.842 ^ 0.197 0.258 ^ 0.151 0.113 ^ 0.022 0.097 ^ 0.119 0.175 ^ 0.070

HgMn Stars with Isotopic Components Resolved 0.004 ^ 0.005 0.001 ^ 0.001 0.0 0.001 ^ 0.001 0.002 ^ 0.003 0.003 ^ 0.003 0.002 ^ 0.001 0.0 0.001 ^ 0.001

11 Per . . . . . . . . . HR 1800 . . . . . . HR 4072 . . . . . . • CrB . . . . . . . . . . 28 Her . . . . . . . . . HR 7245 . . . . . . 46 Aql . . . . . . . . . HR 7775 . . . . . . 74 Aqr . . . . . . . . .

0.075 ^ 0.069 0.001 ^ 0.002 0.0 0.000 ^ 0.002 0.003 ^ 0.003 0.005 ^ 0.010 0.006 ^ 0.006 0.002 ^ 0.002 0.015 ^ 0.015

0.234 ^ 0.024 0.007 ^ 0.002 0.001 ^ 0.001 0.036 ^ 0.004 0.001 ^ 0.001 0.268 ^ 0.019 0.013 ^ 0.010 0.000 ^ 0.001 0.249 ^ 0.019

0.025 ^ 0.042 0.009 ^ 0.002 0.005 ^ 0.001 0.026 ^ 0.005 0.003 ^ 0.003 0.233 ^ 0.014 0.004 ^ 0.001 0.001 ^ 0.001 0.236 ^ 0.023

HgMn Stars with Isotopic Components Not Resolved 0.006 ^ 0.008 0.037 ^ 0.045 0.004 ^ 0.007 0.015 ^ 0.015 0.048 ^ 0.066 0.003 ^ 0.004 0.001 ^ 0.001 0.004 ^ 0.005 0.005 ^ 0.006 0.006 ^ 0.007 0.010 ^ 0.024

i Cnc . . . . . . . . . . n1 Boo . . . . . . . . t Her . . . . . . . . . . / Her . . . . . . . . . HR 6532 . . . . . . 46 Dra A . . . . . . 46 Dra B . . . . . . 112 Her . . . . . . . HR 7361 . . . . . . HR 7664 . . . . . . HR 8349 . . . . . .

196 3.5

198 199

0.026 ^ 0.025 0.042 ^ 0.078 0.047 ^ 0.061 0.014 ^ 0.023 0.063 ^ 0.058 0.002 ^ 0.005 0.002 ^ 0.004 0.006 ^ 0.009 0.038 ^ 0.019 0.002 ^ 0.004 0.053 ^ 0.072

200 202 201

0.156 ^ 0.156 0.073 ^ 0.090 0.025 ^ 0.041 0.014 ^ 0.012 0.049 ^ 0.038 0.001 ^ 0.001 0.004 ^ 0.006 0.015 ^ 0.017 0.206 ^ 0.138 0.044 ^ 0.033 0.014 ^ 0.017

204

46 Aql 28 Her 11 Per 74 Aqr HR 7245

2

HR 7775 HR 1800

1.5

1

0.221 ^ 0.091 0.229 ^ 0.089 0.154 ^ 0.095 0.041 ^ 0.059 0.224 ^ 0.146 0.080 ^ 0.143 0.008 ^ 0.012 0.109 ^ 0.168 0.139 ^ 0.166 0.134 ^ 0.107 0.270 ^ 0.212

weighted toward heavy (short wavelength) or light (long wavelength) isotopes, however. Figure 5 shows the observed and synthetic spectra of • CrB with isotopic and hyperÐne components of j4358 and the included Fe lines indicated as described in the caption. While the isotopic ratios cannot be determined with any certainty, it is not difficult to see that the line is centered near the heavier isotopesÏ components (shorter wavelengths for this line). This is similar to the result found for Hg II using j3984 for • CrB. The Hg I abundances and abundance upper limits are compared with Hg II abundances in Figure 6. The abundances agree fairly well, except for HR 4072 and perhaps • CrB. For HR 4072, A(Hg II) was found to be 0.47 dex larger than A(Hg I), a di†erence larger than the sum of the uncertainties of the abundances. This is similar to the 0.43 dex di†erence found by Smith (1997), although for his work the di†erence was within the joint formal errors for the abundances. For • CrB, A(Hg I) was found to be 0.36 dex larger than A(Hg II), also a di†erence larger than the sum of the uncertainties. SmithÏs A(Hg II) agrees better with both his

3

2.5

0.265 ^ 0.189 0.063 ^ 0.057 0.043 ^ 0.054 0.032 ^ 0.041 0.112 ^ 0.128 0.001 ^ 0.001 0.005 ^ 0.005 0.102 ^ 0.176 0.378 ^ 0.259 0.101 ^ 0.152 0.164 ^ 0.133

HR 4072

TABLE 7 Hg I ABUNDANCES

.5 3983.7

3983.8

3983.9

3984

3984.1

3984.2

FIG. 4.ÈObserved spectra of sharp-lined HgMn stars, ordered by strength of 3984 AŽ feature. Solid vertical lines indicate wavelengths for even Hg isotopic components. Dotted vertical lines indicate wavelength ranges for the hyperÐne components of the two odd Hg isotopes.

Star

A(Hg I)

Hg I Central j (4350 AŽ ])

W j (mAŽ )

11 Per . . . . . . . . HR 1800 . . . . . . l Cnc . . . . . . . . . HR 4072 . . . . . . • CrB . . . . . . . . . . / Her . . . . . . . . . 28 Her . . . . . . . .

[6.05 6.25 ^ 0.10 [5.50 6.17 ^ 0.11 6.10 ^ 0.12 5.89 ^ 0.12 5.59 ^ 0.13

... 8.331 ... 8.329 8.330 ... ...

... 10.5 ... 12.5 7.1 D3.0 D3.0

No. 1, 1999

MERCURY IN HgMn STARS

423 5.

1.05

1

.95 flux

even Hg

.9 Fe lines

.85

.8 4358.1

4358.2

4358.3

4358.4

4358.5

FIG. 5.ÈWavelengths of Hg I isotopic components of j4358 are indicated with respect to a plot of the observed (solid curve) and synthetic (dotted curve) spectrum of • CrB. Vertical lines show locations of even Hg isotopes (solid lines), 199Hg (short-dashed lines) 201Hg (medium-dashed lines), and Fe (dot-dashed lines).

and our A(Hg I). We used m \ 0.9 for • CrB, while Smith used m \ 0.2, which caused much of the di†erence. The central wavelength could be found only for the three sharp-lined stars with fairly strong 4358 AŽ lines. The central wavelengths and the equivalent widths are listed in Table 7. As stated earlier, there is an uncertainty of about 0.005 AŽ in the central wavelength determination that comes from the uncertainty in the wavelength calibration. The central wavelengths listed in Table 7 are those found after setting the Y II line to j4358.728. The central wavelengths of HR 1800, HR 4072, and • CrB are within ^0.002 AŽ of each other, which is small compared with the 0.004 AŽ wavelength di†erence between even Hg I isotopic components for the 4358 AŽ line. In each case Hg I j4358 can be Ðtted acceptably with the Hg isotopic mixture found with Hg II j3984, but a variety of mixtures can be found that Ðt equally well.

6.5

6

DISCUSSION

The cause of the peculiar elemental and isotopic abundances in HgMn and other CP stars has been a puzzle since their discovery. It was suggested soon after their detection that the peculiar abundances in peculiar A and B stars are the result of nuclear and/or accretion processes : they are evolved stars that have returned to the main-sequence region of the H-R diagram, a close binary companion has dumped processed material on them (Fowler et al. 1965 ; Guthrie 1971), or elements are magnetically accreted onto the stars from the interstellar medium (Havnes & Conti 1971 ; Havnes 1974). Such scenarios now seem unlikely. HgMn and other CP stars make up too large a fraction of A and B stars for it to be likely that they all have evolved through the giant phase and are temporarily back near the main sequence on the H-R diagram. Additionally, some very young clusters such as Orion OB1 and Scorpius OB1 contain HgMn stars, which cannot be old enough to have passed through the giant phase. The elemental abundance patterns observed cannot be matched by any plausible nucleosynthesis model. Until recently, radiation pressureÈdriven di†usion, in which photon momentum transferred to absorbing atoms levitates certain elements in the stellar atmosphere or envelope, seemed the most convincing explanation proposed for the peculiar abundances in HgMn stars. Other e†ects such as light-induced drift (LID), disorganized magnetic Ðelds, and selective winds have also been proposed as possibly inÑuencing the abundances. In this paper we have reported Hg elemental abundances for 42 HgMn stars, more than any other single work. We report Hg isotopic abundance estimates for 20 HgMn sharp-lined HgMn stars. Most of these have not been previously analyzed for Hg isotopic composition. Our results agree fairly well with those of other groups for stars that have been previously observed. A number of stellar parameters could conceivably have e†ects on radiative di†usion or other separation mechanisms : T , log g, v sin i, age, the eff the initial and current presence of a binary companion, chemical composition, and so on. In the following sections we will look for correlations between Hg elemental and isotopic abundances with stellar parameters in order to Ðnd clues to the cause of the peculiar abundances observed in HgMn stars. The larger sample of stars and the very high resolution spectra in this study may provide better insights than previous works. We will then discuss di†usion and other possible causes for the peculiar abundances observed in HgMn stars.

A(Hg I)

5.1. E†ects of Stellar Parameters Figure 7 shows3 Hg II abundances versus T for HgMn eff using Hg stars and 66 Eri A. We Ðnd, as Smith (1997) did abundances from ultraviolet spectra, that there is no clear correlation between Hg abundances and T . There seems to be a decrease in A(Hg) with increasing T eff for the sharplined stars (circles). This may be an artifact eff due to the small number of sharp-lined HgMn stars observed. The presence of a binary companion does not seem to a†ect the Hg abun-

5.5

5 5

5.5

6

6.5

A(Hg II)

FIG. 6.ÈHg abundances from Hg II j3984 and Hg I j4358 compared

3 Figures 7, 8, 15, and 16 include data from Smith (1997) for stars not in this paper (87 Psc, / Phe, s Lup, HR 6997, and b Scl), with s Lup indicated as a sharp-lined star and an abundance correction made for the di†erent gf-value used by Smith. SmithÏs data are included to add extra data points to the analysis.

424

WOOLF & LAMBERT 7

6

A(Hg II)

5

4

3

2 11000

12000

13000

14000

15000

16000

FIG. 7.ÈHg II abundances vs. T . Circles are the sharp-lined stars. eff Triangles are upper limits for Hg II for stars with weak 3984 AŽ lines. Squares are all other stars. Open symbols are stars in binaries. Lines connect components of 66 Eri and 46 Dra.

dance. It is interesting that T does not seem to a†ect the eff because T does seem to Hg abundance for HgMn stars, eff play a role in determining whether a star becomes a HgMn star or not. HgMn stars lie in a rather narrow temperature range : 10,500 K [ T [ 16,000 K. Hg II j3984 would be eff strong and easily detectable for, say, main-sequence stars with T \ 8500 or 19,000 K if they had Hg abundances eff observed in HgMn stars, although in cooler commonly stars it may be masked somewhat by the Cr I and Fe I lines listed in Table 3. Preliminary results of a search for HgMn star traits in hot Am stars (at the cool boundary of the HgMn star range) have not found an enhanced j3984 or UV Hg lines enhanced by more than 1.0È1.5 dex (Wahlgren & Dolk 1998). It appears that the abrupt drop in Hg II j3984 line strength at the HgMn star T boundaries is due to drops in Hg abundance, not changeseffin line strength due to, say, Hg going to Hg I at lower temperatures or Hg III at higher temperatures. Figure 8 shows Hg II abundances versus log g for the HgMn stars. It appears that there may be a slight increase of the average A(Hg II) with increasing log g, but it is quite small compared with the scatter in the data and may not be signiÐcant.

Vol. 521

Figure 9 shows A(Hg II) versus v sin i for our data plus s Lup from Smith (1997). One selection criterion for our study was a reported v sin i [ 35 km s~1, so there is a bias in v sin i for our data. Furthermore, there is an observational bias in discovering HgMn stars : it is more difficult to discover a HgMn star with high v sin i and slight overabundances of Mn and Hg than a HgMn star with low v sin i and strong overabundances of Mn and Hg. There may be a slight trend in our data of higher Hg abundance with higher v sin i, but this could be entirely due to selection e†ects. The inset in Figure 10 shows the locations of the HgMn stars studied in this work on the H-R diagram using Hipparcos absolute magnitudes (M ) and the V [I values Hp used in the Hipparcos study, where M \ Hp ] 5 log n Hp [ 10, Hp is the Hipparcos apparent magnitude, and n is the parallax (ESA 1997). No correction was made for binary companions. The line on the plot shows the approximate location of the main sequence from the Hipparcos data ; it is simply a Ðt-by-eye estimate from Figure 3.5.2 in volume 1 of the Hipparcos and Tycho Catalogues (ESA 1997). The outlying points on the plot are HR 1846 at (0.010, [4.67) and HR 1883 at (0.090, [2.62). They were identiÐed as possibly being HgMn stars (Cowley 1972), but we can Ðnd no strong Mn II lines for them in our j/*j \ 60,000 spectra, and as seen in Figure 3a, or a strong Hg II j3984. Their location on the H-R diagram and the fact that they have the two lowest log g values of any of the stars in this study implies that they have evolved o† the main sequence. They are not HgMn stars. In the main panel of Figure 10 the HgMn stars are placed on the H-R diagram with symbols indicating their total Hg II abundance as measured by j3984. Location on the H-R diagram does not predict total Hg abundance for HgMn stars. 5.2. Isotopic Ratios for Sharp-lined Stars Several interesting results come from the analysis of the very high resolution spectra of the sharp-lined HgMn stars (Fig. 4 ; Table 6) obtained using the 2.7 m telescopeÏs coude spectrometer at McDonald Observatory. One notable result is that all of the sharp-lined stars have essentially no 196Hg or 198Hg. Whatever e†ect is producing the high Hg abundances, at least in the sharp-lined stars, is not concen-

7 6.5 6 6

5 A(Hg II)

A(Hg II)

5.5

5

4 4.5 3

4

3.5

2 3.2

3.4

3.6

3.8

4

4.2

FIG. 8.ÈHg II abundances vs. log g. Symbols are the same as in Fig. 7. Solid lines connect components of 66 Eri and 46 Dra.

0

10

20

30

40

50

FIG. 9.ÈHg II abundance vs. v sin i. Open circles indicate upper limits on v sin i, and inverted triangle indicates upper limit on A(Hg II).

No. 1, 1999

MERCURY IN HgMn STARS -5

-2 0

5

-1 -.5

0

.5 V - I [mag]

1

1.5

0

1

2 -.3

-.2

-.1

0

.1

.2

V - I [mag]

FIG. 10.ÈHgMn stars on H-R diagram with symbols indicating Hg II abundances. Open squares indicate stars with A(Hg) ¹ 4.5, open triangles indicate stars with 4.5 \ A(Hg) \ 6.0, and asterisks indicate stars with A(Hg) º 6.0. The inset shows location of stars from this study on the Hipparcos H-R diagram. No correction was made for binary companions ; M is the Hipparcos absolute magnitude. Hp

trating the two lightest Hg isotopes in the atmospheric region where j3984 is formed. This may not be surprising for 196Hg, which is a trace isotope making up 0.15% of the solar system Hg mixture. However, 198Hg makes up 10.0% of the solar system Hg mixture, so its absence in HgMn stars is signiÐcant. Although the elemental abundances observed in HgMn stars do not appear to be caused by nuclear processes, it could be signiÐcant that 196Hg and 198Hg are the two stable Hg isotopes that are not produced by the r-process ; the b-decay path to them is blocked by the stable 196Pt and 198Pt isotopes. The other stable Hg isotopes can be r-process products. The trace isotope 196Hg is produced by the p-process. A few other recent studies (Wahlgren & Dolk 1998 ; Proffitt et al. 1999 ; Jomaron, Dworetsky, & Bohlender 1998 ; Smith 1997) report Hg II isotopic mixtures in HR 7775 that cannot be adequately described by the ““ q-parameter ÏÏ introduced by White et al. (1976). The q-parameter is a dimensionless mix parameter that describes abundances of Hg isotopes relative to 202Hg through the following two equations : log a \ q(A [ 202) log

10

e,

(1)

[N(A)/N(202)] * , (2) [N(A)/N(202)] _ where N(A) is the nonlogarithmic abundance of the isotope with atomic mass A, and e is the base for natural logarithms. The ratio of the abundance of isotope A to that of 202Hg is increased or decreased by the factor a relative to its terrestrial abundance ratio. Our spectrum of the HR 7775 Hg II j3984 gives similar results. Two other sharplined stars observed in this study, 11 Per and HR 7245, also have 3984 AŽ lines that cannot be well described using q. This is the Ðrst time that the very unusual Hg isotopic ratios found in 11 Per and HR 7245 have been reported. The Hg II isotopic mix of HR 7245 shows approximately equal abundances of four isotopes with small or zero abundances for the other three : 0% 196Hg, 0% 198Hg, 27% 199Hg, 23% 200Hg, 4% 201Hg, 24% 202Hg, and 21% 204Hg. Such a pattern is difficult to explain using radiative di†ua\

425

sion. It may be that the isotopic ratios are set by the Hg being concentrated until the lines are saturated, which would help explain the similar abundances for 199Hg, 200Hg, 202Hg, and 204Hg, but this would not explain the absence of 198Hg or the small 201Hg abundance. Michaud, Reeves, & Charland (1974) described a Ðnely balanced di†usion process in which the small mass di†erences between Hg isotopes are sufficient to cause isotope separation in HgMn star atmospheres. If this type of di†usion is occurring in HgMn stars, then the substantial di†erence in radiation force caused by hyperÐne splitting may be enough to push 201Hg out of the line-forming region, leaving other isotopes behind. In HR 7775 the spectrum shows about 37% 202Hg and 62% 204Hg. An adjustment of the stellar parameters used to calculate the Hg abundance, especially the microturbulence, would change the calculated 202Hg/204Hg ratio somewhat, but it would not explain the absence of the lighter isotopes. The failure of the q-parameter to describe the Hg isotopic mix of HR 7775 (or other HgMn stars) requires that any model, di†usion or otherwise, have forces so Ðnely balanced that it can cause a sharp cuto† in concentrations of the Hg isotopes at 202Hg, pushing the lighter elements o† of HR 7775 or to a level where they are not observable as Hg II. The pattern seen in 11 Per is even more difficult to explain. Its Hg II as observed in j3984 is mostly 199Hg and 204Hg. There were no detectable changes in the Hg isotope ratio for 11 Per, which was observed Ðve times (1994 December, 1995 February and December, 1997 September, and 1998 December), and no other absorption lines appear doubled in any of the spectra. Thus it is unlikely that the strange isotopic ratio determined here is due to a binary companion contributing one of the two main components seen in j3984. It is difficult to explain why di†usion would concentrate the heaviest Hg isotope and a lighter isotope without concentrating the intermediate isotopes. It is possible that the shorter wavelength component of j3984 belongs to some other element, perhaps a twice-ionized species, which would only have such a strong line in a hotter star like 11 Per (T \ 14,549 K). But during the decades eff no one has identiÐed such a blendj3984 has been studied, ing line. It does not seem probable that such a blending line would be absent in all other HgMn stars. We do not believe it likely that the short-wavelength component is due to another element, although we can think of no other explanation for the Hg isotopic abundance pattern seen for 11 Per. HR 1800, • CrB, and HR 4072 are very similar HgMn stars. Their Hg II isotopic abundances are similar. Their temperatures are in the range 11,000 ^ 500 K, their log g values fall between 3.65 and 3.80, their microturbulences are less than 1 km s~1, they have very small v sin iÏs, and each is in a binary. It is difficult to believe that interaction with the binary companion is causing the peculiar abundances, however. HR 1800 is in a visual binary with a period of 338.5 yr (Starikova 1983), and • CrB and HR 4072 are SB2Ïs with periods of 35.47 and 11.58 days, respectively (Dworetsky 1980 ; Nariai 1976). All three show 201Hg in their j3984 spectra. Whatever is acting to remove 201Hg from HR 7245 is not doing so for these three stars. The Hg in 46 Aql and 28 Her is almost pure 204Hg, as it is in s Lup, which was not included in this study because its declination ([33¡) makes it difficult to observe from McDonald Observatory. The physical parameters of s Lup

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and 28 Her are very similar. For 28 Her, T \ 10,908, eff log g \ 3.83, and m \ 0.5. For s Lup, T \ 10,560, eff log g \ 3.90, and m \ 0.0. However, 28 Her appears to be single and has v sin i \ 9 km s~1, while s Lup is in an SB2 with a period of 15.2 days (Batten, Fletcher, & MacCarthy 1989) and has v sin i \ 1 km s~1. The HgMn star 46 Aql is di†erent from the other two ““ pure 204Hg ÏÏ HgMn stars. It has T B 12,900 K, about the middle of the range of temeff for HgMn stars, and log g \ 3.74. peratures There is no glaring di†erence between the parameters of the pure 204Hg group and the group including HR 1800, • CrB, HR 4072, and HR 7775, which have Hg II made up of mostly 202Hg and 204Hg with a small amount of the lighter isotopes. They are all cool and have log gÏs that are about average for HgMn stars. If some parameter or combination of parameters can be found that discriminates between these two groups, it may provide a strong lead in determining the cause of the isotopic shuffling that occurs in HgMn stars. The Hg isotopic mix of 74 Aqr, which has a Hg elemental abundance 4.0 dex larger than the solar system abundance, is the closest to solar of any of the sharp-lined HgMn stars observed here. Figure 11 compares the spectrum of 74 Aqr with a synthetic spectrum calculated using the solar Hg isotopic mixture and the Hg II abundance found for 74 Aqr. In the Ðgure we can see that even 74 Aqr is depleted in 198Hg and enhanced in 204Hg compared with the solar mixture. It may be that other stars have unusual Hg II isotopic mixes similar to those of these sharp-lined stars, but the other stars observed here have larger v sin iÏs or were not observed at very high resolution, so their isotopic ratios cannot be unambiguously determined. With the exception of HR 2676, which has a very weak 3984 AŽ line, none of the HgMn stars observed for this work have 3984 AŽ line central wavelengths shorter than that of 74 Aqr (3983.933 AŽ ), which implies that the higher v sin i stars also have Hg isotopic mixtures that are at least slightly enhanced in the heavy isotopes and may be missing 196Hg and 198Hg as well. For sharp-lined stars where estimates of isotopic abundances are possible, we can directly compare the 204Hg abundances (Table 9). This can be done for other isotopes as well, but 204Hg is present in all HgMn stars that

1

Vol. 521

6.5

6

5.5

5

4.5

3.5

3.6

3.7

3.8

3.9

4

4.1

4.2

FIG. 12.È204Hg II abundance vs. log g for sharp-lined HgMn stars. Filled circles indicate stars with isotopic components resolved. Open circles indicate stars with unresolved isotopic components.

show enhanced Hg abundances, and the isotopic abundance ratio of 204Hg to other Hg isotopes is always enhanced relative to the solar system mixture, so it is especially of interest. In Figure 12 we see that there is no obvious trend in A(204Hg) with log g when the values are plotted for the sharp-lined stars from this work and s Lup, except that we see no stars with both log g \ 3.7 and A(204Hg) [ 5.5. In Figure 13 A(204Hg) is plotted against T for the sharp-lined stars. We saw in Figure 7 that A(Hg) eff decreases with T for the sharp-lined stars but not for the eff the observed HgMn stars observed for combination of all this work, and in Figure 16 we will see that the fraction of Hg in heavy isotopes increases with decreasing T for the combination of all the observed stars. Thus it is eff not clear whether the decrease in A(204Hg) with T for the very sharp lined stars is due to a decrease in totaleffHg abundance or in the fraction of heavy Hg isotopes. The trend is not present for the sharp-lined stars as a whole. If the higher v sin i stars are actually acting di†erently than the sharplined stars, as indicated in Figure 7, then v sin i needs to be studied in more depth for its e†ect on abundances in HgMn

6.5

.95 6

.9 5.5

.85

.8

5

.75 4.5

3983.7

3983.8

3983.9

3984

3984.1

3984.2 10500

FIG. 11.ÈObserved 74 Aqr spectrum (solid line) compared with synthetic spectrum with solar system meteoritic Hg isotopic mixture (dashed line).

11000

11500

12000

12500

13000

13500

14000

14500

FIG. 13.È204Hg abundance vs. T for sharp-lined HgMn stars. eff Symbols are the same as in Fig. 12.

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MERCURY IN HgMn STARS

427

4.05

4.05

4

4

3.95

3.95

3.9

0

10

20

30

40

50

11000

12000

13000

14000

15000

16000

FIG. 14.ÈCentral wavelength of Hg II j3984 vs. v sin i. Open circles indicate upper limits on v sin i.

FIG. 16.ÈCentral wavelength of j3984 vs. T . Symbols are the same as eff in Fig. 15.

stars and its implications in models for producing the peculiar abundances. The central wavelength of the Hg II j3984 is another indication of the isotopic mixture of Hg in HgMn stars. It can provide isotopic information for broader lined stars in which the isotopic mixture estimates using synthetic spectra are not very precise. Longer wavelengths correspond to heavier Hg isotopes (Preston 1973 ; Cowley & Aikman 1975). In Figure 14 we plot the central wavelength of the 3984 AŽ feature versus v sin i. The three stars with almost pure 204Hg as indicated by a central wavelength near 3984.072 AŽ all have v sin i \ 10 km s~1. This is probably signiÐcant. The bias introduced by selecting low v sin i stars for observation and the possibility of larger errors in Ðnding the central wavelength for higher v sin i stars must still be taken into account, but if broader lines are causing errors in measuring the central wavelength, they are not scattering any up to the pure 204Hg wavelength. We do not know what the low v sin i, pure 204Hg correlation means for models used to explain HgMn star peculiar abundances. Figures 15 and 16 plot the central wavelength of j3984 versus log g and T , respectively. No signiÐcant correlation eff the Hg isotopic mixture and log g. The can be seen between

isotopic blend is correlated with T . The stars with mostly eff with the exception of heavy Hg isotopes tend to be cooler, 46 Aql, which has almost pure 204Hg and T \ 12,914 K. The only star with a central wavelength foreffj3984 shorter than 3983.93 AŽ , the wavelength expected for a solar system meteoritic isotope ratio, is HR 2676, which has a weak 3984 AŽ line that may be strongly a†ected by noise in the observed spectrum. If we ignore this star, it seems that all HgMn stars have Hg II isotope ratios that are about solar or weighted more toward the heavier isotopes, which may give some clue as to what is causing the high Hg abundances. Presence of a binary companion does not seem to a†ect the isotopic mix. The average central wavelength of the 3984 AŽ feature for HgMn stars in binaries in this work is 3.983, while the central wavelength for single HgMn stars is 3.978, an insigniÐcant di†erence. There is no signiÐcant trend in Hg abundance versus the central wavelength of j3984. It is possible to group the HgMn stars based on the apparent Hg isotope mixtures (Table 8). ClassiÐcations were simple for sharp-lined stars but were less certain for stars with higher v sin i. The groups correspond closely to central wavelengths of j3984. Stars with very broad or very weak lines were not classiÐed ; s Lup is included in the classiÐcation. Groups are deÐned as follows : (1) Pure 204Hg, j º 3984.05 AŽ ; (2) mostly 204Hg and 202Hg, 3984.02 AŽ \ j \ 3984.05 AŽ ; (3) concentrated toward heavy Hg isotopes but includes some lighter ones, 3983.97 AŽ \ j \ 3984.02 AŽ ; (4) Hg not enhanced in heavy isotopes, j ¹ 4983.97 AŽ ; and (5) 11 Per, the oddball. It is possible that there are other stars with very strange isotope mixtures like that of 11 Per that have been missed because rotational broadening has prevented their detection. In Figure 17 the upper section of the H-R diagram, where HgMn stars lie, is shown with the symbols indicating the starsÏ Hg isotopic mix groups. As expected from previous work, the stars with mostly heavy Hg isotopes (circles and squares) are at the cool end of the diagram. No pattern is seen for stars without the heavy Hg isotopes enhanced (asterisks).

6.5

A(Hg II)

6

5.5

5

4.5

4 3.9

3.92

3.94

3.96

3.98

4

4.02

4.04

4.06

4.08

FIG. 15.ÈCentral wavelength of j3984 vs. log g. Circles are the 10 sharp-lined stars. Squares are other stars. Open symbols indicate the star is in a binary. The components of 46 Dra are connected by a line.

5.3. Possible Causes of Peculiar Abundances 5.3.1. Di†usion For the past 25 years or so, the most commonly accepted explanation for the peculiar elemental and isotopic abun-

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WOOLF & LAMBERT

Vol. 521

TABLE 8

TABLE 9

STARS GROUPED BY LINE SHAPE OF j3984

204Hg ABUNDANCES FOR SHARP-LINED STARS

Star

Groupa

Star

A(204Hg)

23 Cas . . . . . . . . . 11 Per . . . . . . . . . 53 Tau . . . . . . . . 66 Eri A . . . . . . . 66 Eri B . . . . . . . k Lep . . . . . . . . . . HR 1800 . . . . . . HR 1846 . . . . . . HR 1883 . . . . . . 33 Gem . . . . . . . c CMa . . . . . . . . . HR 2676 . . . . . . HR 2844 . . . . . . f CMi . . . . . . . . . HR 3273 . . . . . . HR 3361 . . . . . . 14 Hya . . . . . . . . l Cnc . . . . . . . . . . i Cnc . . . . . . . . . . HR 4072 . . . . . . HR 4493 . . . . . . c Crv . . . . . . . . . . HR 4817 . . . . . . n1 Boo . . . . . . . . s Lup . . . . . . . . . . • CrB . . . . . . . . . . t Her . . . . . . . . . . HR 5998 . . . . . . / Her . . . . . . . . . 28 Her . . . . . . . . . HR 6532 . . . . . . 38 Dra . . . . . . . . HR 7018 . . . . . . HR 7028 . . . . . . 46 Dra A . . . . . . 46 Dra B . . . . . . 112 Her . . . . . . . HR 7245 . . . . . . HR 7361 . . . . . . 46 Aql . . . . . . . . . HR 7664 . . . . . . HR 7775 . . . . . . 30 Cap . . . . . . . . HR 8349 . . . . . . 74 Aqr . . . . . . . . . 69 Peg . . . . . . . . .

4 5 X X 3 4 3 X X 4 X X 4 X X 2 X 2 4 3 X 3 4 4 1 3 3 X 3 1 4 4 3 3 3 2 3 4 4 1 4 2 X 4 4 X

Sun . . . . . . . . . . . . 11 Per . . . . . . . . HR 1800 . . . . . . HR 4072 . . . . . . • CrB . . . . . . . . . . 28 Her . . . . . . . . HR 7245 . . . . . . 46 Aql . . . . . . . . HR 7775 . . . . . . 74 Aqr . . . . . . . .

0.006 4.87 5.93 6.43 5.29 5.29 4.45 4.75 6.23 4.24

direction to support the observed overabundance of Hg and could plausibly produce the Hg isotopic anomalies in HgMn stars. In one version of their model, radiative forces in the photosphere push Hg II ions up to the level in the atmosphere where they are ionized to Hg III because of the lower electron pressure there. The radiative forces on Hg III are much smaller. So Hg is concentrated there with the radiation and gravitational forces approximately equal and N(Hg II) D N(Hg III). The close balance between radiation and gravitational forces may allow the small mass di†erence between isotopes to produce abundance anomalies. Di†usion may be acting in HgMn stars. If there is no outer convection zone or other mixing mechanism near the surface of HgMn stars, it would be difficult to explain why it would not be acting in them. It may be possible to explain the bias toward heavy Hg isotopes in some HgMn stars by a di†usion process, perhaps with a mechanism similar to that described in Michaud et al. (1974) where di†usion forces are Ðnely balanced and the isotopic mass di†erences are sufficient to cause the abundance di†erences. More detailed calculations are needed, but it may be possible to use di†usion to explain the approximately equal abundances of 199Hg, 200Hg, 202Hg, and 204Hg and the lack of 201Hg in HR 7245 if abundances increase until the lines are saturated, but the hyperÐne splitting of 201Hg does

-2

-1

a Groups are deÐned as follows : (1) Pure 204Hg ; (2) Mostly 204Hg and 202Hg ; (3) Concentrated toward heavy Hg isotopes but include some lighter ones ; (4) Hg not enhanced in heavy isotopes ; (5) 11 Per, the oddball ; and (X) the line was too weak or too broad to classify.

dances observed in HgMn stars has been that radiation pressureÈdriven di†usion alters the initial chemical composition by concentrating and depleting the various elements in the starsÏ outer layers. Michaud (1970) suggested that radiative di†usion may cause the peculiar abundances in CP stars. Then Michaud et al. (1974) showed that di†usive forces were of the right magnitude and in the right

0

1

2 -.3

-.2

-.1

0

.1

.2

V - I [mag]

FIG. 17.ÈHgMn stars on H-R diagram with symbols indicating Hg isotope mixtures determined from the Hg II j3984. Symbols are as follows : stars with pure 204Hg (group 1 ; open circles), mostly 204Hg and 202Hg (group 2 ; open squares), concentrated heavy isotopes with some light isotopes (group 3 ; open triangles), no concentration of heavy isotopes (group 4 ; asterisks), and 11 Per (group 5 ; Ðlled pentagon).

No. 1, 1999

MERCURY IN HgMn STARS

not allow its line to saturate, so that the isotope is pushed out of the line-forming region. If the 1% mass di†erence between 202Hg and 204Hg is sufficient to allow di†usion to create pure 204Hg stars, then a radiative force on 201Hg that is more than 20% larger than the forces on other Hg isotopes could easily cause a low 201Hg abundance like that seen in HR 7245. But it is not obvious what causes the di†erences between stars with similar parameters. If hyperÐne splitting causes depletion of 201Hg in HR 7245, then why does it not do so in HR 1800, • CrB, or HR 4072 ? There seems to be no way to use di†usion to explain the bizarre Hg isotopic mix found in 11 Per. It may be possible to explain the missing 201Hg in HR 7245 by hyperÐne splitting, but 11 Per has low abundances of 200Hg, 201Hg, and 202Hg, sandwiched between high abundances of 199Hg and 204Hg. Mass di†erences cannot be invoked to explain this isotopic mix. The even isotopes do not have hyperÐne splitting to desaturate their lines and increase radiative acceleration. Di†usion fails as an explanation. There now seem to be some major problems with di†usion-only models in addition to the Hg isotopic anomalies just discussed. Proffitt et al. (1999) have done calculations that show that the radiative force on Hg ions in s Lup is not stronger than the gravitational force on them at any layer in the photosphere : the Hg abundance observed in s Lup cannot be supported by radiative forces alone. 5.3.2. E†ect of HyperÐne Splitting on Di†usive Forces

The absence of 201Hg in HR 7245 while 200Hg and 202Hg are present points to the possibility that hyperÐne splitting is a†ecting the concentration of Hg in the line-forming region. To test the e†ects of hyperÐne splitting of Hg lines on di†usive forces, calculations were made of radiative acceleration on Hg isotopes using the modiÐed version of I. HubenyÏs SYNSPEC version 42u FORTRAN code, provided by C. R. Proffitt (1998, private communication). The code assumes LTE for its atomic line analysis. The input model atmosphere was the same LTE Kurucz model used for the abundance analysis of HR 7245. HR 7245 was used because it is the star that highlighted the question of hyperÐne splittingÏs e†ect. The principle behind the modiÐcations to SYNSPEC is the same as the modiÐcations made to do non-LTE radiation force calculations using I. HubenyÏs TLUSTY program as described in Proffitt et al. (1999). The modiÐcations to the SYNSPEC code did not change the basic operation of the code, except to add calculations of radiative force. For each wavelength it starts at the bottom layer of the atmosphere, gets the necessary parameters (T , eff opacities, density, and so on) by reading them in from the model stellar atmosphere or calculating them, determines which lines in the line list contribute opacity at that wavelength, and determines what fraction of the photons at that wavelength get passed up to the next layer in the atmosphere. This procedure continues up to the top layer of the atmosphere, each layer starting with the photon Ñux passed to it from the layer below it. Radiative force is calculated at each level by Ðnding the fraction of the photon Ñux absorbed by a species, say Hg II, and Ðnding what that translates to in force per ion of that species. The Hg line list, also provided by C. R. Proffitt, included atomic data calculated for Proffitt et al. (1999) and reported by Brage, Proffitt, & Leckrone (1999). The calculations used the model atmosphere for HR 7245 and all the Hg lines

429

between 750 and 13000 AŽ . In one run, the line list used was unchanged. In the second, the hyperÐne components of 201Hg were combined into one isotopic component for each line. No lines of other elements were included. Additional calculations were run with and without hyperÐne splitting for individual Hg lines. HyperÐne splitting makes little di†erence for acceleration on Hg I. This is not surprising, since the Hg I lines are weak, and desaturation of the lines by hyperÐne splitting should be a minor e†ect. When hyperÐne splitting is included, the acceleration on 201Hg II due to radiation pressure increases by 20%È75% in deeper layers where the acceleration is strong and by up to 400% in layers near the surface where desaturation due to hyperÐne splitting lets more light through in the model. In the layers where radiative acceleration is strongest, including hyperÐne splitting increases the radiative acceleration on 201Hg III by about 20%È55%. HyperÐne splitting was included for a limited number of lines in the original Hg line list, but it was included for the stronger lines. The di†erence in the derived radiative accelerations would be somewhat larger if a more complete line list were used for calculations with and without hyperÐne splitting. SYNSPEC was also run on small wavelength intervals to test hyperÐne splittingÏs e†ect on the radiative acceleration due to individual Hg lines. As could be expected, in relatively weak lines (e.g., Hg I j4358) and lines with large isotope splitting (e.g., Hg II j3984), where there is less line saturation to begin with, the e†ect was small, less than 1%. For stronger lines with smaller isotope splitting, the e†ect was larger. Including hyperÐne splitting increased the calculated radiative acceleration by 20%È60% in atmospheric layers where acceleration is strong in Hg III j1647 and by 200%È900% in Hg II j1942. Isotopic and hyperÐne component wavelengths for jj1647 and 1942 are listed in Table 3 of Proffitt et al. (1999). The code used for the calculations assume LTE. Using the non-LTE code would probably give di†erent increases in the calculated radiative force because of the inclusion of hyperÐne splitting, but these LTE calculations are sufficient to show that hyperÐne splitting has a large e†ect on the radiative force on 201Hg. 5.3.3. Microturbulence and Di†usion

Di†usion processes such as molecular di†usion (di†usion due to the random thermal motions of the particles) and thermal di†usion (di†usion due to the vertical temperature gradients in the stellar atmospheres), which could conceivably undo abundance concentrations because of radiation pressureÈdriven di†usion, have been included in the di†usion calculations of the various authors. It is not clear that the possibility of mixing due to microturbulence has been considered. Microturbulence is still not well understood. As used in stellar elemental abundance calculations, it is deÐned as nonthermal motions of absorbing particles along the line of sight on a length scale that is small compared with the line-forming region. If microturbulent motions are acoustic waves that propagate without a†ecting the relative arrangement of particles in the atmosphere, then microturbulence probably does not result in mixing sufficient to undo radiative di†usion e†ects. If microturbulent motions in upper main-sequence stars are small-scale, e†ectively random mass motions, then it may be that there is mixing

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WOOLF & LAMBERT

due to microturbulence. The efficiency of mixing attributable to such random motions depends strongly on their mean free path. The root mean square (rms) displacement, d, of a particle in a ““ random walk ÏÏ is approximately equal to the geometric mean of its mean free path, l, and its total distance traveled, D : d B JlD \ Jlv6 t ,

(3)

where v6 is the rms velocity of the particle and t is time. For such displacements to not cause mixing sufficient to undo radiative di†usion e†ects, we need l[

d2 . v6 t

(4)

If we assume a di†usion timescale of about 104 yr (Michaud 1970) and a microturbulent velocity of about 1 km s~1, and we want displacements smaller than a line-forming region on the order of hundreds of kilometers, then we get l[

(100 km)2 \ 3 ] 10~5 m . (1 km s~1)(10,000 yr)

(5)

Measurements of microturbulence only require that the scale of the motions are small compared with the lineforming region. Scales of meters, centimeter, or microns would Ðt observations equally well. If, as seems likely, microturbulence is mass motions on scales larger than 10~4 m, then it may undo the e†ects of radiative di†usion. If microturbulence is acoustic waves or if it is mass motions on scales of microns, then it probably does not. A better understanding of the nature of microturbulence is needed to determine which case holds. 5.3.4. L ID

LID has also been suggested as a possible cause for the peculiar abundances in HgMn stars. It occurs when there is an asymmetric light Ñux distribution across the Dopplerbroadened absorption proÐle of an atomic line. The asymmetry causes the absorbing particles to be selectively excited based on their velocities (Atutov & Shalagin 1988). If the upper and lower atomic states of that particular line have di†erent collisional cross sections, then they are decelerated di†erently by the gas at rest, and there is a net Ñux of absorbing particles. The isotope splitting of a line itself may provide the Ñux asymmetry needed to drive LID and cause isotopic separation. Nasyrov & Shalagin (1993) showed that particle separation due to LID can be much stronger than the radiative di†usion e†ects. However, LeBlanc & Michaud (1993) argued that since Hg is mostly ionized in HgMn stars and the collisional cross sections of ionized species are not as dependent on the excitation state, LID should have a negligible e†ect on Hg in HgMn stars. Proffitt et al. (1999) propose a way in which LID might be able to work without neutral species, but more detailed calculations are needed to determine if it will work. A further problem with LID as the mechanism producing the overabundances observed in HgMn stars is that the well-studied HgMn star s Lup has enhanced abundances of almost all the elements observed heavier than iron. The light Ñux asymmetries needed for LID are presumably produced by strong absorption lines of other elements, which should be e†ectively randomly distributed with respect to the LID a†ected species. An element with a line on the ““ blue ÏÏ wing of another line will feel a net acceleration

Vol. 521

toward the center of the star, while if the elementÏs line is on the ““ red ÏÏ wing of another line, it will feel a net acceleration toward the surface of the star (Nasyrov & Shalagin 1993). It is difficult to believe that all the heavy elements happen to have strong absorption lines on the proper side of other elementsÏ absorption lines to cause overabundances. Alternatively, if isotope splitting provides the needed Ñux asymmetries for LID and the splitting usually splits the components of heavier isotopes to longer wavelengths, then that may partly explain why most of the heavier elements are enhanced. Then, however, one would have to explain why elements like Au and Bi, with only one stable isotope and thus no isotope splitting, are overabundant. 5.3.5. Other E†ects

With the measurement of a 3.6 kG quadratic magnetic Ðeld in the HgMn star 74 Aqr (Mathys & Hubrig 1995), it becomes important to consider the e†ects of magnetic Ðelds on the abundances in HgMn stars. It is possible that other HgMn stars also have complicated magnetic Ðelds that are difficult to detect observationally but that a†ect abundances. Michaud, Megessier, & Charland (1981) showed that magnetic Ðelds can strongly a†ect radiative di†usion. Observations to determine if other HgMn stars have nondipole magnetic Ðelds should be performed. If a sizable fraction of HgMn stars have (previously unobserved) magnetic Ðelds, then it may be that magnetic Ðelds are playing a major role in producing the abundance and isotopic anomalies observed in HgMn stars. It may be that several mechanisms are acting together to produce the observed abundances in HgMn stars. Examples are the suggestions of Proffitt et al. (1999) that a combination of radiative forces on Hg IV below the photosphere and weak mixing of Hg up into the photosphere produce a concentration of Hg, that LID and selective wind loss may then produce the Hg isotopic mixtures we see, or that an isotope separation e†ect based on isotopic mass like that suggested by Michaud et al. (1974) may be working below the photosphere where Hg IV goes to Hg III and producing the isotopic mixture we observe. This is similar to the case suggested for Tl (Proffitt, Brage, & Leckrone 1996). Tl could not be supported in the atmosphere of s Lup by radiative forces, but there are sufficient radiative forces just below the photosphere to support it. More detailed quantitative calculations are needed to determine whether the model is capable of producing the observed overabundances and isotope anomalies for Hg. The fact that we Ðnd that the Hg abundance in HgMn stars is not correlated to T while the abundances of Mn eff (Smith & Dworetsky 1993b), and a few other elements are and that there exist HgMn stars enhanced in Mn but not Hg (e.g., 53 Tau) or enhanced in Hg but not Mn (e.g., s Lup), is evidence that there is more than one e†ect causing the peculiar abundances or that, if there is a single cause, it a†ects some elements very di†erently. For any model explaining the peculiar elemental and isotopic abundances observed in HgMn stars to be fully accepted, it is important that the combination of mechanisms used in the model work for every element observed or not observed in the stars. For example, if ProffittÏs suggestion of Hg IV di†usion below the photosphere and weak mixing into the photosphere (Proffitt et al. 1999) can be shown to explain the observed Hg abundances, then it should also produce the abundances observed for other ele-

No. 1, 1999

MERCURY IN HgMn STARS

ments or at least not destroy the concentrations and depletions observed if they are produced by another mechanism. Furthermore, to be fully accepted, a model used to explain the peculiar elemental and isotopic abundances in HgMn stars must explain the similarities and di†erences between HgMn stars, other CP stars, and normal stars of similar temperatures, gravities, rotational velocities, and so on. If the model cannot produce the di†erent abundances present in a HgMn star, a He weak star, and a normal star with similar parameters, then it is incomplete. It seems that no mechanism has yet been proposed that is able to explain fully the elemental and isotopic abundances present in HgMn stars. Further work is needed to determine if some combination of proposed mechanisms and, perhaps, some previously unexamined mechanisms can account for them. Future investigations should include searches for nondipole magnetic Ðelds in additional HgMn stars, further theoretical calculations of element concentration models

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that include multiple mechanisms and many atomic elements and stars, non-LTE abundance analyses for key elements, further studies of HgMn stars in binaries and clusters, studies of Hg abundances in stars near the HgMn star temperature boundaries, and a correlation of results for di†erent classes of CP stars. Thanks to C. Proffitt for providing suggestions and the radiative force code, M. Dworetsky and R. Napiwotzki for providing help with the code to Ðnd log g and T using eff uvby b photometric data, U. Litzen for obtaining FTS spectra of Hg II j3984, C. Sneden for help with MOOG, and G. Michaud and J. Scalo for comments helpful in producing this work. This research has made use of the SIMBAD database, operated at the Centre de Donnees Astronomique de Strasbourg, France. The authors acknowledge the support of the National Science Foundation (grant AST 96-18414) and the Robert A. Welch Foundation of Houston, Texas.

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