first direct evidence that barium dwarfs have white dwarf ... - IOPscience

The Astronomical Journal, 141:160 (8pp), 2011 May C 2011.

doi:10.1088/0004-6256/141/5/160

The American Astronomical Society. All rights reserved. Printed in the U.S.A.

FIRST DIRECT EVIDENCE THAT BARIUM DWARFS HAVE WHITE DWARF COMPANIONS R. O. Gray1 , C. E. McGahee1 , R. E. M. Griffin2 , and C. J. Corbally3 1

Department of Physics and Astronomy, Appalachian State University, Boone, NC 28608, USA; [email protected], [email protected] 2 Herzberg Institute of Astrophysics, National Research Council, 5071 West Saanich Road, Victoria, BC V9E 2E7, Canada; [email protected] 3 Vatican Observatory Research Group, Steward Observatory, Tucson, AZ 85721-0065, USA; [email protected] Received 2010 May 25; accepted 2011 February 21; published 2011 April 4

ABSTRACT Barium ii (Ba) stars are chemically peculiar F-, G-, and K-type objects that show enhanced abundances of s-process elements. Since s-process nucleosynthesis is unlikely to take place in stars prior to the advanced asymptotic giant branch (AGB) stage, the prevailing hypothesis is that each present Ba star was contaminated by an AGB companion which is now a white dwarf (WD). Unless the initial mass ratio of such a binary was fairly close to unity, the receiving star is thus at least as likely to be a dwarf as a giant. So although most known Ba stars appear to be giants, the hypothesis requires that Ba dwarfs be comparably plentiful and moreover that they should all have WD companions. However, despite dedicated searches with the IUE satellite, no WD companions have been directly detected to date among the classical Ba dwarfs, even though some 90% of those stars are spectroscopic binaries, so the contamination hypothesis is therefore presently in some jeopardy. In this paper, we analyze recent deep, near-UV and far-UV Galaxy Evolution Explorer (GALEX) exposures of four of the brightest of the class (HD 2454, 15360, 26367, and 221531), together with archived GALEX data for two newly recognized Ba dwarfs: HD 34654 and HD 114520 (which also prove to be spectroscopic binaries). The GALEX observations of the Ba dwarfs as a group show a significant far-UV excess compared to a control sample of normal F-type dwarfs. We suggest that this ensemble far-UV excess constitutes the first direct evidence that Ba dwarfs have WD companions. Key words: binaries: spectroscopic – stars: chemically peculiar – stars: evolution – stars: individual (HD 2454, HD 15306, HD 26367, HD 34654, HD 114520, HD 221531) – white dwarfs few Ba dwarfs (in particular relative to the numbers of the more massive Ba giants) have been identified, and whether those Ba dwarfs have properties consistent with them being progenitors of Ba giants. Intrigued by this unresolved mystery, we have started to investigate the identification and properties of Ba dwarfs. The identification of new members of the class is especially important, as much of what we know about these rare stars comes from statistical studies. Because of the small numbers of known Ba dwarfs, those statistical results (including the ones reported in this paper) are necessarily tentative. That study is still ongoing, and we report here what we believe are the first direct detections of WD companions of Ba dwarfs.

1. INTRODUCTION The spectral peculiarities associated with barium ii (Ba) stars, in particular their enhanced abundances of s-process elements, invoke a choice between two basic explanations: internal or external. However, the choice has not seemed a difficult one to make: s-process nucleosynthesis is unlikely before the star has reached the asymptotic giant branch (AGB) stage of evolution, and since most Ba giants were shown by McClure et al. (1980) to be members of binary systems and at least some proved to have enhanced UV fluxes (see, for instance, Böhm-Vitense et al. 2000) suggestive of white dwarf (WD) secondaries, the explanation of external contamination of the present Ba star by a companion during its relatively rapid terminal evolution to a WD has appeared the more likely of the two. Processes of diffusion are also unlikely causes of the Ba star abundance phenomena on account of the existence of deep surface convection zones and the long diffusion timescale for these stars (on the order of the lifetime of late-F and early-G dwarfs with Teff 6300 K; Vauclair & Vauclair 1982). Possibly through their initial recognition as a sub-group of cool giants, Ba stars have tended to become synonymous with Ba giants; the paper by McClure (1983) assumes so unquestioningly, as indeed do many others, but unless the initial mass ratio of the binary system in question was close to unity, then the star that is contaminated by its evolving component is at least as likely to be a dwarf as a giant. Even so, the sample of Ba dwarfs known to date is considerably smaller than that of Ba giants. Simple arguments based on evolutionary timescales and the orbital characteristics of Ba dwarfs and Ba giants (North et al. 2000) on the one hand, and WD cooling times (Böhm-Vitense et al. 2000) on the other hand, suggest that most Ba giants must have been contaminated while on the main sequence, implying that the Ba dwarfs are progenitors of Ba giants. The question is therefore raised as to why so

2. THE SUB-GROUP OF Ba DWARFS The known barium dwarfs are main-sequence or slightly evolved stars with spectral types between F4 and early G. The first examples of stars that would later become classified as Ba dwarfs were recognized by Bidelman (1981) who noticed, on objective-prism plates, a number of F and early-G dwarfs with unusually strong Sr ii lines. The Ba dwarfs are characterized by a significant overabundance of the s-process elements, in particular strontium and barium (Edvardsson et al. 1993; North et al. 1994; Gray & Griffin 2007). We have to recall that the magnetic Ap (Sr) stars also exhibit enhanced strontium; however, those stars are believed to acquire their peculiar abundances via diffusion and chemical separation, whereas the Ba dwarfs are generally much cooler than Ap (Sr) stars and have deep surface convection zones, resulting in such long diffusion timescales that the chemical peculiarities of the Ba dwarfs cannot be explained by diffusion. Nor can they be old enough or luminous enough to be in a stage of evolution in which s-process elements have been synthesized in advanced nuclear processes; that does not happen until a star has reached the 1

The Astronomical Journal, 141:160 (8pp), 2011 May

Gray et al.

Table 1 GALEX Photometry, Spectral Types, and Physical Parameters of the Ba Dwarfs Star HD 2454 15306 26367 34654 114520 221531

GALEX Fluxes NUVb

Sp. Type

Dist.a

V

MV

Mass

Teff

log g

[M/H]

Binarity

FUV (mag)

(Continued)

(pc)

(mag)

(mag)

(M )

(K)

(mag) 11.620 0.001 12.631 0.001 11.947 0.000 12.494 0.003 12.742 0.003 12.470 0.000

15.174 0.001 16.855 0.007 7.354 0.008 8.629 0.076 5.480 0.018 7.364 0.006

F5 V Sr ii

36.6 0.4 ... ... 38 2 60 7 141 14 96 10

6.04

3.24 0.03 ... ... 3.68 0.11 3.39 0.25 1.07g 0.22 3.43 0.23

1.45

6500

4.04

−0.32

?

1.71

6660

3.83

−0.29

RV varc

1.30

6220

4.15

+0.07

SB1d , ABe

1.12

6070

4.35

−0.09

RV varf , ABe

1.62

6570

3.90

+0.02

SB2f , ABe

1.35

6510

4.18

−0.11

RV varc , ABe

F4 V Sr ii (kF1mF2) F7 V+ Sr ii (CH+0.4) F8.5 V Sr ii F5 IV-V Sr ii F5 V Sr ii

8.92 6.56 7.29 6.86 8.33

Notes. Uncertainties in the GALEX magnitudes, the parallaxes, and the absolute magnitudes are indicated below the quantities. a Based on the revised Hipparcos parallaxes (van Leeuwen 2007). b The GALEX NUV magnitudes for all the stars in this table lie within the nonlinear regime of the GALEX detector and are therefore unreliable. c North & Duquennoy (1991). d Gray & Griffin (2007). e Astrometric Binary, Makarov & Kaplan (2005). f This paper. g Combined M for the visual and spectroscopic systems. V

ratio (S/N) = 100 in the stellar continuum in that band. We also located archived GALEX observations of sufficient accuracy for two more stars which we have recently recognized as Ba dwarfs (HD 34654 and HD 114520), bringing our sample size to six. Table 1 lists those six Ba dwarfs along with the final GALEX NUV and FUV photometry. Classification-resolution spectra (1.8 Å/2 pixels) were obtained with the 0.8 m telescope at Appalachian State University’s Dark Sky Observatory (DSO), located in the Blue Ridge mountains of North Carolina, G/M Spectrograph and a 1200 line mm−1 grating. The spectra were recorded on a thinned, back-illuminated 1024 × 1024 Tektronix CCD operating in the multipinned-phase mode; they extend from 3800–4600 Å. An Fe–Ar hollow-cathode comparison lamp was observed for wavelength calibrations, and the spectra were reduced with standard IRAF4 routines. Our new, precise spectral types are included in Table 1. All six stars are of spectral type F. The two new Ba dwarfs mentioned above had been discovered with the DSO equipment on the basis of their strong Sr ii λλ 4077 and 4215 Å lines. Those discovery spectra are illustrated in Figure 1 along with the spectra of two MK standards. HD 114520 is a visual binary with ΔV ≈ 4.0 and a separation of about 10 . It was also discovered to be a spectroscopic binary by Parsons (1983). HD 34654 is an astrometric binary (Makarov & Kaplan 2005). High-resolution spectra of HD 34654 and HD 114520 were then obtained with the coudé spectrograph in highdispersion mode (96 inch (2.4 m) camera, mosaic grating with 830 line mm−1 , and a Richardson image slicer) at the coudé focus of the 1.2 m telescope of the Dominion Astrophysical Observatory (DAO) in Victoria, Canada. Those spectra, recorded on an SITe-4 CCD, have a dispersion near 2.4 Å mm−1 and cover a wavelength range of ∼140 Å with a resolving power ∼45,000 in the blue and ∼30,000 in the red. The wavelength

AGB. The contamination hypothesis of McClure et al. (1980) therefore seems a likely explanation for both the Ba giants and the Ba dwarfs. There is, however, a significant problem with the McClure hypothesis to explain the existence of Ba dwarfs. On the positive side, North et al. (2000) pointed out that most known Ba dwarfs are single-lined spectroscopic binaries. They were able to show statistically that the unseen companions in those binaries have a mean mass of 0.67 M , and we have indeed found that the mass of the companion of the Ba dwarf HD 26367 is 0.60 M (Gray & Griffin 2007). Masses of that magnitude are close to what one would expect for WDs arising from AGB progenitors of mass 1.5–3.5 M (Catalán et al. 2009; Dominguez et al. 1999). However, on the negative side, searches by North & Lanz (1991) in IUE data—and Bond (1984) made a similar search for a related class, the CH subgiants—have failed to turn up any direct evidence that the companions of Ba dwarfs are WDs. There are other problems with the hypothesis that Ba dwarfs are the progenitors of the Ba giants which are beyond the scope of this paper to address. One of the most important is that the known Ba dwarfs all have evolutionary masses 1.7 M (see Table 1), whereas Ba giants may have masses as high as 5 M (Böhm-Vitense et al. 2000). In this paper, we pursue the observation (North et al. 2000) that Ba dwarfs appear to be in binary systems, but confine our investigation to the ultraviolet where any evidence of hot companions should be fairly unequivocal. 3. OBSERVATIONS We obtained time on the Galaxy Evolution Explorer space telescope (GALEX; Martin et al. 2005), and report here observations, completed in 2009 October, of the four brightest Ba dwarfs (HD 2454, HD 15306, HD 26367, and HD 221531) in the far-UV (FUV) and near-UV (NUV) bands. The flux from a companion with Teff 10,000 K, even if a WD, would dominate the flux in the FUV band (λeff = 1539 Å) in all four stars, so exposure times were calculated so as to reach a signal-to-noise

4

IRAF is distributed by the National Optical Astronomy Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation.

2

Gray et al. 12.5

Ba II 4554

Mg II 4481

Y II Fe I

Sr II 4216

Ca I 4226

G-band

Hδ

Fe II, Ti II

Fe I 4046

Sr II 4077

Ca II H

Ca II K


12.0

Hγ

Control Sample F-type Dwarfs Barium Dwarfs

11.5

Rectified Intensity

HR 5634 F5 V

11.0 10.5

*

10.0

HD 114520 F5 IV-V Sr II

*

F-V

0

9.5 9.0

*

HD 34654 F8.5 V Sr II

*

8.5 8.0 7.5

η Cas A F9 V

7.0 3900

4000

4100

4200

4300

4400

4500

6.5 0.20

4600

Wavelength ( )

0.22

0.24

0.26

0.28

0.30

0.32

0.34

0.36

0.38

b-y

Figure 1. Classification spectra of the two newly recognized Ba dwarfs, HD 114520 and HD 34654, compared to two MK standards, HR 5634 and η Cas A. Note the strong Sr ii lines (marked with ∗) in the Ba dwarfs.

Figure 2. F − V colors, formed from the GALEX FUV magnitude (F) and the Johnson-V magnitude, plotted against Strömgren b − y. The control sample of F dwarfs is represented by open circles, and the Ba dwarfs by filled stars. Average errors for the control-sample F dwarfs are σ (F − V ) ≈ 0.05 mag and σ (b − y) ≈ 0.005 mag. A regression line through the control sample is shown for reference. The effect of reddening is indicated by the arrow.

calibration source was an Fe–Ar hollow-cathode lamp for the blue spectra and a Th–Ar lamp for the red ones. Observations in the blue were centered near 3933, 4101, 4150, and 4481 Å, and in the red near 6170 Å. The spectra were reduced with a semi-automatic IRAF-based pipeline and were then corrected for 3.5% scattered light as recommended by Gulliver et al. (1996) and Adelman et al. (2006).

Before the positions of the Ba dwarfs can be interpreted, it is necessary to discuss some of the general properties of this color–color plot. The relationship between F − V and b − y for the control sample appears to be nearly linear, although its scatter of σ ∼ 0.25 mag is considerably larger than expected on the basis of the average uncertainty of ∼0.05 mag in the values of F − V. That extra scatter may arise from the following sources:

4. ANALYSIS 4.1. The Color–Color Diagram

1. Interstellar extinction, which is much greater in the FUV band than in the optical. However, because the F − V scale is compressed relative to that of the b − y axis in Figure 2, reddening will displace a star along a diagonal line, as indicated by the arrow. The slope of that reddening line has been calculated from the standard reddening law (Cardelli et al. 1989). Even though the reddening line roughly parallels the regression line, if the amount of reddening is large it could contribute substantially to the scatter in the diagram. As a matter of fact, the stars in the control sample are all within 40 pc, so their reddening should be negligible. We verified that this was the case by using Crawford’s method (Crawford & Mandwewala 1976) to calculate the E(b−y) color excess from Strömgren uvbyβ photometry for each individual star. We found that the resulting values of E(B−V ) = E(b−y) /0.74 mag are distributed normally about E(B−V ) = −0.001 ± 0.014 mag, as would be consistent with negligible reddening for the entire group. However, four of the six Ba dwarfs are more distant than 40 pc, so their reddenings must be examined individually (Section 4.2). 2. Metallicity, which has an important effect on the FUV flux (in the sense that metal-weak stars will show significant FUV excesses) since most of the opacity—both continuous and line—in the FUV bandpass is due to metals. Many Ba dwarfs are slightly metal weak (North et al. 1994), so we have been careful to choose our control sample such that it has a range of metallicities that spans that of the Ba dwarfs (see Figure 3). We determined the basic physical parameters (Teff , log g, and [M/H], where [M/H] is the measure of the overall metallicity) from DSO 1.8 Å resolution spectra for the Ba dwarfs and the control sample with the “simplex method” of Gray et al. (2001, 2003)

We commenced with the observation of North et al. (2000) that most known Ba dwarfs are single-lined binaries, and tailored our analysis so as to investigate the properties of the companion stars. If a companion has Teff 10,000 K and is a WD, its flux will dominate the stellar photospheric flux in the GALEX FUV band for even the hottest known Ba dwarf (F4; Teff ≈ 6700 K). However, ascertaining the existence of a far-ultraviolet excess is not straightforward. Theoretical fluxes in the FUV band for F-type stars are not well determined because of uncertainties in both the continuous and line opacities. Chromospheric emission can also increase the flux in the FUV band. The best approach is therefore an empirical one that compares the Ba-dwarf fluxes with those of F-type dwarfs of similar Teff , surface gravities, metallicities, ages, and chromospheric activity. Because two of our Ba dwarfs unfortunately have poorly determined distances, and one (HD 15306) does not even have a measured parallax, comparison of absolute fluxes is not practical. Instead, we employed a “color–color diagram” to make the flux comparisons. For the dependent variable, we had hoped to use a purely ultraviolet color, namely FUV–NUV, but it proved difficult to find F-type dwarfs with both archived GALEX observations of sufficient accuracy in the FUV-band and NUV-band observations that were neither saturated nor defective in some other way. Instead, we adopted the FUV–VJohnson color (hereinafter symbolized F − V) for one axis of the color–color diagram and the Strömgren b − y index for the other. From the NStars database (Gray et al. 2003, 2006), we extracted 58 stars with spectral types F1–F9 and which have archived GALEX observations such that σ (FUV) < 0.10 mag; they are referred to hereinafter as the “control sample.” In Figure 2 they are plotted in the [F − V, b − y] plane along with the six Ba dwarfs. 3


Gray et al.

16

12

Control Sample Ba Dwarfs

14

11 12

[M/H] = 0.0

Number

F-V

10

8

10

9

[M/H] = -0.5 6

8 4

7 0.20

2

0.22

0.24

0.26

0.28

0.30

0.32

0.34

0.36

0.38

b-y 0 -0.5

-0.4

-0.3

-0.2

-0.1

0.0

0.1

Figure 5. Color–color diagram of Figure 2 with theoretical lines based on ATLAS9/SPECTRUM (Castelli & Kurucz 2003; Gray & Corbally 1994) fluxes for two different metallicities, [M/H] = 0.0 and −0.5. The zero point for the theoretical F − V colors has been chosen such that the [M/H] = 0.0 curve aligns approximately to the upper envelope of the control-sample distribution. The same zero point has been applied to the [M/H] = −0.5 colors.

0.2

[M/H]

Figure 3. Comparing the Ba dwarfs to the control sample of F dwarfs, in terms of [M/H]. Black bars indicate Ba dwarfs and gray ones the control sample. 8.75

Table 2 Correlation Results

9.00

3.4 9.25

Correlation Indicator

3.6

[M/H] [Fe/H] [α/H]

9.50

3.8

log(g)

2.0 9.75

4.0 1.7

1.25

7000

98% 95% 97%

all those of the Ba dwarf sample: [M/H], Teff , log g, and age (Figure 4). To understand better the effects of metallicity on the F − V index, we modeled the photometry in the GALEX FUV passband (Morrissey et al. 2007) corresponding to theoretical stellar energy distributions derived from the new ATLAS9 models for two metallicities, [M/H] = 0.0 (solar metallicity) and −0.5, and added them to Figure 2 (see Figure 5). The line for [M/H] = −0.5 is well separated from that for [M/H] = 0.0, suggesting that metallicity differences can indeed give rise to scatter in the controlsample observations. In fact, the values of [M/H] for the stars in the control sample and their vertical deviations from the regression line in Figure 2 [Δ(F − V )] show a correlation that is highly significant (Pearson r = 0.502, for which the null hypothesis may be rejected with a confidence >99.99%). It is of interest, however, to investigate this correlation further. As it turns out, most of the important sources of opacity in the FUV band in F dwarfs arise from α-capture and light odd-Z, even-N elements (such as Al) and not iron-peak elements. To see if we could derive a better correlation with Δ(F − V ) using the abundance of α elements, we measured both [Fe/H] and [α/H] (as represented by the mean of [Ca/H] and [Si/H] calculated from equivalent widths of four Ca i and nine Si i lines in the 6100–6250 Å region) in the 15 control-sample stars with measurable Elodie spectra. For those 15 stars, we derive the correlation results shown in Table 2. As anticipated, the [α/H] abundances give a slightly stronger and more significant correlation with the Δ(F −V )

4.6 7200

0.585 0.507 0.574

1.0

ZAMS

4.4

7400

Significancea

Note. a Defined as 100% × (1—the probability of the null hypothesis).

1.5

4.2

Pearson Correlation Coefficient

6800

6600

6400

6200

6000

5800

Teff

Figure 4. Showing the positions of the control sample of F dwarfs (open circles) and Ba dwarfs (filled stars) in the theoretical H-R diagram. Also included are solar-abundance isochrones (solid lines labeled with log(ages)) and evolutionary tracks (dashed lines, labeled with masses in units of solar masses). The isochrones and evolutionary tracks are from Lejeune & Schaerer (2001). This figure demonstrates that the control sample of F dwarfs spans the Ba dwarfs considered in this paper in terms of effective temperature, log g, and age. The evolutionary masses of the Ba dwarfs may also be estimated from this figure (see Table 1). The cross indicates the position of a “mean” Ba dwarf, in terms of Teff and log g.

but employing models computed with the new version of ATLAS9 (Castelli & Kurucz 2003). Unlike the models computed with the “old” opacity distribution functions of Kurucz (1993), the new models incorporate more modern (and therefore presumably more accurate) abundances.5 The physical parameters of the control sample thus span 5

As part of a final omnibus publication on the NStars program (described by Gray et al. 2003, 2006), we have checked the new simplex [M/H] values by computing [Fe/H] values (employing Teff and log g from the simplex fits and the Fe i lines used in Gray & Griffin 2007) for a sample of 38 F- and G-type stars from high-resolution spectra downloaded from the Elodie archive (Moultaka et al. 2004). The simplex [M/H] computed with the new ATLAS9 models are in close agreement with [Fe/H] values based on the high-resolution spectra and, taking into account the uncertainties in the [Fe/H] values, show an intrinsic scatter of ∼0.07 dex and a zero-point offset of 0.02 dex. (For more details, see R. O. Gray et al. 2011, in preparation.)

4


Gray et al.

correlation arises from star-to-star variations in the calcium (and by extension α-capture element) abundances in the controlsample stars. We can test that possibility statistically in the set of 15 control-sample stars for which we have measured α-element abundances (see above). In that set, ΔS2 proves to be uncorrelated with either [α/H] (−0.105; 29%) or [Fe/H] (−0.113; 31%). In the whole control sample, there is no significant correlation between ΔS2 and the simplex [M/H] values (−0.089; 50%), and so it seems safe to assume that ΔS2 and [M/H] are independent variables and that the correlation between ΔS2 and Δ(F − V ) is unaffected by metallicity effects.

0.14

0.13

S2

0.12

0.11

0.10

0.09

4.2. Notes on Individual Ba Dwarfs

0.08 0.20

0.22

0.24

0.26

0.28

0.30

0.32

0.34

0.36

HD 2454 = HR 107. The brightest and first identified Ba dwarf (Tomkin et al. 1989), HD 2454 is one of the four target stars for which we obtained high-S/N GALEX photometry. Located at a distance of 36.6 pc (van Leeuwen 2007), its reddening, determined using Crawford’s method, is near zero (E(B−V ) ∼ −0.003 ± 0.014); the uncertainty was estimated from the spread of the control-sample reddenings. A simplex fit (Gray et al. 2003) to a classification-resolution spectrum and Strömgren uvby photometry yields Teff = 6578 K, log g = 4.45, and [M/H] = −0.20, in close agreement with its spectral type (Table 1).

0.38

(b-y)

Figure 6. Chromospheric activity index S2 as a function of (b − y). S2 rises steadily with increasing temperature (decreasing b − y). We use a different index, ΔS2 , defined as the vertical displacement of each measured S2 value from the lower envelope of the distribution (solid line), in order to compensate for that trend.

values than do the [Fe/H] abundances, but interestingly the [M/H] values from the simplex method perform equally well. This can probably be ascribed to the fact that the [M/H] values are derived from the metallic-line spectrum in the 3800–4600 Å region, which contains a mixture of both iron-peak and α elements, including calcium, magnesium, and titanium. 3. Chromospheric activity, which—even if low—can give rise to excess FUV emission, displacing points vertically downward in Figure 2. Since chromospheric activity is a strong function of age in F- and G-type stars (Soderblom et al. 1991; Mamajek & Hillenbrand 2008), we were careful to choose stars for the control sample which had ages spanning those of the Ba dwarfs (see Figure 4). To investigate the effects of chromospheric activity on Δ(F − V ) in detail, we defined a chromospheric-activity index S2 (measured on the DSO 1.8 Å resolution spectra), which is closely related to the Mount Wilson S-index (Vaughan et al. 1978) but determined with a 2 Å wide passband in the Ca ii H and K cores instead of a 1 Å wide passband (for more details, see R. O. Gray et al. 2011, in preparation), and measured it in both the control sample and the Ba dwarfs. Repeated observations of chromospherically quiet (i.e., non-variable) stars indicated a typical measuring uncertainty in S2 of ∼ ±0.001. Our S2 indices for the control sample are plotted against b − y in Figure 6.

HD 15306. This star was first classified as “F; str λ4077,” by Bidelman (1985) from objective-prism plates taken for the University of Michigan Southern Spectral Survey. Some of those “F; str λ4077” stars were later placed in the Ba dwarf category by North & Duquennoy (1991) on the basis of an abundance analysis. North & Duquennoy (1991) examined HD 15306 with a number of techniques and deduced a low value for the reddening (E(B−V ) ≈ 0.01), a distance of about 150 pc, variable radial velocity (RV, indicating binarity), and a sub-solar metallicity ([Fe/H] = −0.34 ± 0.09). HD 15306 is another of the four Ba dwarfs for which we obtained high-S/N GALEX photometry. The interpretation of the GALEX data for HD 15306 is particularly problematic because (1) it has no Hipparcos parallax, so we do not know its distance or its reddening, apart from the estimates of North & Duquennoy and (2) no Strömgren photometry is available. It has, however, been observed in the Geneva photometric system (Rufener & Maeder 1971) and that photometry may be reliably transformed to Strömgren photometry. We have derived our own transformations from Strömgren and Geneva observations of the control sample of F dwarfs as well as of other stars of similar spectral types chosen from the NStars database. The transformations are valid for F0–G5 dwarfs with [M/H] > −0.5:

Figure 6 indicates a relationship that clearly climbs in the direction of bluer colors (smaller b − y), an effect that can be attributed to increasing amounts of photospheric flux in the 2 Å cores of the Ca ii H and K lines with increasing effective temperature. To remove that trend, we defined a new index, ΔS2 , as the vertical displacement between a measured S2 and the lower envelope as drawn in that figure; we then attributed any remaining scatter to varying amounts of activity in the control sample. In the control sample, ΔS2 shows a robust and significant correlation (−0.468; 99.98%) with the vertical deviations from the regression line in Figure 2 [Δ(F − V )]. That correlation is anticipated, because chromospheric emission in the cores of the Ca ii H and K lines should certainly be strongly related to emission in the chromospheric and transition-region lines in the FUV region. However, it is possible that part of that

y=V b−y v−b u−v

= = = =

0.010 + 0.998Vm 0.132 + 0.712(B2 − V1 ) 0.829 + 1.000(B1 − B2 ) 0.639 + 1.029(U − B1 )

σ σ σ σ

= 0.032 = 0.006 = 0.010 = 0.021

where the Strömgren quantities are on the left, and the Geneva magnitudes and indices are on the right. For HD 15306, those equations yield y = V = 8.930, b − y = 0.261, m1 = 0.154, and c1 = 0.590. The “foreground” galactic extinction along the line of sight to HD 15306, quoted by NED and calculated via the extinction maps of Schlegel et al. (1998), is E(B−V ) = 0.033 mag. This color excess should be regarded as an upper limit to the actual E(B−V ) for HD 15306, so it would appear that HD 15306 is in 5


Gray et al. ◦

a low-reddening part of the sky (Galactic latitude = −54. 9). Unfortunately, no measurements of the Hβ index are present in the literature, so it is not possible to derive a value for the reddening by Crawford’s method, though an estimate may also be obtained by comparing the spectral type (F4 V) with entries in Table 2 of Gray et al. (2001). For F4 V stars, the mean (b − y)0 color is given in that table as 0.271 ± 0.013 mag, which is redder than the b − y value for HD 15306, suggesting that for HD 15306, E(B−V ) = E(b−y) /0.74 ∼ −0.014 ± 0.019 mag; the uncertainty is based on the values given in Table 2 of Gray et al. (2001). A simplex fit to our classification spectrum of HD 15306 and the above Strömgren indices yields Teff = 6660 K, log g = 3.83, and [M/H] = −0.29 ± 0.07 in reasonably good accord with the results of North & Duquennoy (1991) and the spectral type we give in Table 1. HD 26367. This star, located at a distance of 38 pc, was first reported as a Ba dwarf by Gray & Griffin (2007). They assumed E(B−V ) = 0.00 mag and showed that it is a singlelined spectroscopic binary. No Hβ photometry is available for this star, but a comparison of its b − y color (0.343) with the mean color for spectral type F7 given in Gray et al. (2001) yields an estimated E(B−V ) ∼ 0.027 ± 0.019 mag, where the uncertainty is derived from the data in Table 2 of Gray et al. (2001). Gray & Griffin (2007) went on to estimate the mass of the companion as 0.60 ± 0.06 M . A more definitive orbital solution and a more complete abundance analysis is currently in preparation. HD 34654. This star was identified as a Ba dwarf in the course of a spectroscopic survey at DSO of Hipparcos F-type astrometric binaries (Makarov & Kaplan 2005). Its classification spectrum is illustrated in Figure 1. No Hβ photometry is available. The distance to HD 34654, calculated from the parallax given by van Leeuwen (2007), is 60 pc; from its (b − y) color and spectral type, we estimate E(B−V ) ∼ 0.014 ± 0.019 mag. Spectroscopy obtained at the DAO indicates that this star is an RV variable; other than that variability, no sign of the companion is visible in the spectrum. As mentioned above, Makarov & Kaplan have determined that HD 34654 has an unseen astrometric companion, although no orbital period could be determined. HD 114520. This star is also located at a high galactic latitude (+82.◦ 6); the “foreground” reddening from Schlegel et al. (1998) suggests an upper limit for E(B−V ) of 0.030 mag. Strömgren uvbyβ photometry is available (b–y = 0.271 mag), and Crawford’s method gives E(b−y) ∼ 0.010 ± 0.010 mag. Comparison with Table 2 of Gray et al. (2001) indicates that the observed b − y is actually slightly bluer than the mean (b − y)0 for F5 IV–V stars, implying E(B−V ) = E(b−y) /0.74 ≈ −0.014 ± 0.013 mag. Combining the two estimates for E(B−V ) suggests E(B−V ) ∼ −0.001 ± 0.009 mag. Its distance is 141 pc. The unusual strength of the Sr ii λ4077 line of HD 114520 was first noted by Gray et al. (2001) as a peculiarity, although the star was not identified as a Ba dwarf in that paper. However, a re-examination of the classification spectrum (see Figure 1) revealed the hallmark signatures of a Ba dwarf. To solidify that identification, we obtained high-resolution spectra of HD 114520 at the DAO for abundance analysis and RV information. Parsons (1983) noted that the star is an RV variable and estimated an orbital period of several hundred days. Our RV measurements to date combined with those of Parsons yield a preliminary orbit

solution with a period on the order of 400 days. Our spectra also indicate that the visible component is itself a double-lined spectroscopic binary, the companion being a late-F or early-G dwarf. This system thus offers a rare opportunity to test the contamination hypothesis, since a previous AGB star should have contaminated both stars during its advanced evolution. However, the long orbital period means that it will take time to sort everything out. HD 221531. Bidelman (1981) classified this star as “F; str λ4077” because of its very strong Sr ii λ4077 line. It was later designated as a Ba dwarf by North et al. (1994) on the basis of overabundances of s-process elements. North & Duquennoy (1991) discovered that HD 221531 is a single-lined spectroscopic binary with a period in excess of 800 days. HD 221531 is located at high Galactic latitude (−66◦ ) and thus has low reddening; the extinction maps of Schlegel et al. (1998) suggest an upper limit of E(B−V ) = 0.033 mag. Geneva photometry is available, but no Strömgren photometry. Our Geneva-to-Strömgren transformations yield Strömgren indices of V = y = 8.321, b − y = 0.294, m1 = 0.149, and c1 = 0.473. Its distance, calculated from the Hipparcos parallax (van Leeuwen 2007), is 96 pc. If we compare the mean (b − y)0 for its spectral type (Gray et al. 2001) to the observed one, we derive E(B−V ) = 0.013 ± 0.019. 4.3. Ba Dwarf Masses The evolutionary masses of the Ba dwarfs in our sample, as well as those of the control-sample F dwarfs, may be estimated from the solar-abundance (Z = 0.02) evolutionary tracks (Lejeune & Schaerer 2001) displayed in Figure 4 and listed in Table 1. From that diagram it can be seen that the most massive Ba dwarf in our sample (which appears to be the most massive known) has M ∼ 1.7 M . All of the Ba dwarfs in our sample happen to be F-type stars, but there are G-type Ba dwarfs as well, so the minimum mass is somewhat smaller than 1 M . The masses of the control-sample F dwarfs clearly span those of our Ba dwarfs. 4.4. The Color Excess of the Ba Dwarf Ensemble We wish to investigate whether the Ba dwarfs as a group show a significant FUV excess relative to the control sample of F dwarfs, once effects due to metallicity and chromospheric activity are accounted for. We reported in Section 4.1 that, for the control sample, E¯ (B−V ) = −0.001 ± 0.002 mag (standard error of the mean). The mean reddening calculated from the color excesses reported in Section 4.2 for the six Ba dwarfs is E¯ (B−V ) = 0.003 ± 0.002 mag. Those mean color excesses are not significantly different beyond the 1σ level, so correcting the colors of the Ba dwarfs for reddening is not warranted. Applying that small correction would not in any case materially affect the following discussion and conclusions. The physical parameters recorded in Table 1 were calculated by the simplex method (Gray et al. 2003) on the assumption of zero reddening for each star. 4.5. The FUV Excess of the Ba Dwarf Ensemble F-type stars may show an FUV excess if they are metal weak or are chromospherically active, so we seek to derive a quantity E(F −V ) , which represents the FUV excess of an F-type star corrected for those effects. We have already shown in Section 4.1 that [M/H] and ΔS2 are uncorrelated. Similarly, it turns out that, for the control sample, ΔS2 and (b − y) are 6


Gray et al.

16 14

Control Sample F-dwarfs Ba Dwarfs

0.0

12

-0.5

10

-1.0

9560 K 10250 K

EFUV

Number

12050 K

8

-1.5

6 -2.0

13800 K 14650 K

4 -2.5

16500 K

2 0 -0.8

-3.0 6200

-0.6

-0.4

-0.2

0.0

0.2

0.4

0.6

6250

6300

6350

6400

6450

6500

6550

6600

Ba Dwarf Effective Temperature

0.8

Figure 7. Histograms of the far-ultraviolet excess E(F −V ) for the control-sample F dwarfs (gray bars) and the Ba dwarfs (black bars). Negative values of E(F −V ) indicate excess flux in the FUV band. The Ba dwarfs have a mean E(F −V ) value that is significantly different from that of the control-sample F dwarfs (see the text).

Figure 8. Isothermal contours in the FUV excess (E(F −V ) ) vs. Ba dwarf Teff diagram based upon IUE fluxes of θ Boo and Procyon and the scaled fluxes for ∼0.6 M WDs with known effective temperatures. The single point with the error bars represents the location of the FUV excess of the “model Ba dwarf” defined in Section 5. This allows us to estimate a “typical” effective temperature for the WD companions of the Ba dwarfs considered in this paper—about 11,000 K.

uncorrelated (r = −0.043, 34%) and [M/H] and (b − y) are uncorrelated (r = 0.101, 55%), so those three variables are statistically independent. We therefore employed a linear regression analysis to derive corrections to the FUV excess for metallicity and chromospheric-activity effects in the controlsample and to remove the trend with (b − y). That analysis yielded an expression for E(F −V ) :

significant result. We therefore conclude that the ensemble FUV excess observed in the population of Ba dwarfs is real. One obvious interpretation of that excess is that the Ba dwarfs have hot companions and—since those companions must also have low luminosities—that they are WDs. However, because of the remaining variance in these distributions, none of the individual Ba dwarf excesses should be taken as significant.

E(F-V) = Far Ultraviolet Excess

E(F −V )

=

(F − V )

−

−

0.84[M/H] ±0.18

−

27.13(b − y) ±0.52 1.75 ±0.02,

+

16.56ΔS2 ±3.83

5. DISCUSSION Because we cannot assign statistical significance to the individual FUV excesses for any of the Ba dwarfs in our sample, it is impossible to determine the individual temperatures of the WD companions for which we now find statistical evidence. However, for future reference it is useful to estimate a “typical” temperature for those companions. To do that we first model a “typical” F star Ba dwarf as one that is characterized by the average values of the Teff , gravity, and metallicity of the sample: Teff = 6424 ± 92 K, log g = 4.08 ± 0.09, and [M/H] = −0.14 ± 0.05, where the errors quoted are the standard errors of the means. We then calculate the Teff of a 0.6 M WD which would yield the ensemble Ba dwarf FUV excess, EF −V = −0.25 ± 0.08 mag (standard error of the mean) when its FUV flux is added to the FUV flux of the model Ba dwarf. To do so, we require the FUV flux of a star with the physical parameters of the model Ba dwarf, and such a star does not exist in the IUE archives. However, Procyon and θ Boo have similar surface gravities (log g = 4.05 and 4.12) and metallicities ([M/H] = −0.06 and −0.08, respectively, as determined by simplex fits) and they bracket the model Ba dwarf in Teff ; the IUE archives contain a number of high-S/N short-wavelength spectra of both stars. We therefore combined their IUE fluxes with the scaled IUE fluxes of ∼0.6 M WDs of known Teff (Holberg et al. 2003, we adopted the ones in their Figure 6) and compared them to the FUV excesses (relative to the star without a WD companion) computed via synthetic photometry. The results are illustrated in Figure 8; it suggests that the “typical” low-luminosity companions to the Ba dwarfs in our sample have Teff ≈11000 ± 750 K, where the error should be interpreted as the standard error of the mean.

where the variance in the relationship is σ = 0.18 mag. That scatter is approximately what is expected from the propagation of errors, as using σF −V ≈ 0.05 mag, σb−y ≈ 0.005 mag, σΔS2 ≈ 0.001, and σ[M/H] ≈ 0.07 dex yields an expected σE(F −V ) ≈ 0.21 mag, suggesting that we have accounted for all the important effects producing an FUV excess in these (single) stars. For a given star, a negative value of E(F −V ) indicates excess flux in the GALEX FUV band relative to the mean of the control-sample F dwarfs, after correcting for metallicity and chromospheric activity. Figure 7 shows a histogram of the distribution of E(F −V ) for both the control sample and the six Ba dwarfs. The E(F −V ) distribution for the control sample is, by design, centered on E(F −V ) = 0, while that for the Ba dwarfs appears to be shifted toward more negative values. In fact, we find E(F −V ) = −0.25 mag, with σ = 0.20 mag. The variances of these two populations are therefore similar, so a Student’s t-test can help us decide if the difference in the means is statistically significant. That test yields t = 3.15, with the result that, despite the small sample size of Ba dwarfs, the probability that these two distributions are drawn from populations of different means exceeds 99.8%. An alternative statistical test which may be applied in this situation is the Kolmogorov–Smirnov test, which can assign a probability to the question of whether two distributions are different. It yielded a value for the Kolmogorov–Smirnov statistic of 0.678, with a probability >99.4% that the two samples are drawn from populations with different distributions—again, a highly 7


Gray et al.

Our finding statistical evidence that Ba dwarfs have WD companions lends support to the McClure contamination hypothesis for the creation of Ba dwarfs, but an outstanding question that remains is the relationship between the Ba dwarfs and the Ba giants. As noted in the introduction, both North et al. (2000) (on the basis of evolutionary timescales and orbital characteristics) and Böhm-Vitense et al. (2000) (on the basis of WD cooling times) have argued that most Ba giants must have been contaminated while on the main sequence, thus implying that the Ba dwarfs are progenitors of Ba giants. If that is so in every instance, then, for a given Ba-star mass, the WD companions of Ba dwarfs should be statistically younger and therefore hotter than those of Ba giants. Demonstrating that would help to clarify the relationship between Ba dwarfs and giants. However, the picture is not so simple. First, the present paper deals with the most massive Ba dwarfs known, and all of them have M 1.7 M . The Ba giants investigated by Böhm-Vitense et al. (2000) all have masses 2.0 M , and thus the effective temperatures of the WD companions of the two sets cannot be compared. That is because the main-sequence lifetimes of the Ba dwarfs in our sample are considerably longer than the mainsequence lifetimes of the Ba giants studied by Böhm-Vitense et al. (2000), and thus the WD companions of our Ba dwarfs could have had a longer time to cool than those of the giants. To make a direct comparison, it will either be necessary to detect the companions of lower-mass Ba giants or to discover more massive Ba dwarfs. Another complication is that if the mass ratio of the original binary was near unity, then a Ba star could have been contaminated as a giant. The WD companion of such a Ba giant could be very young (and hot) indeed (see, for instance, the case of WeBo 1, a young Ba giant at the center of a planetary nebula; Bond et al. 2003). Other questions include (1) will all Ba dwarfs evolve to become Ba giants or will mixing in the envelope sufficiently dilute the s-process overabundances so that the star appears normal by the time it becomes a giant and (2) if Ba dwarfs are the progenitors of Ba giants, then where are the Ba dwarfs with masses >1.7 M ? If such massive Ba dwarfs do exist, they would be early-F or A-type stars showing chemical peculiarities. It could then be surmised that some Ap or Am stars may actually be Ba dwarfs; an investigation into that possibility represents another project.

R.O.G., C.E.M., and R.E.M.G. acknowledge support through NASA grant NNX09AD64G under the GALEX Guest Investigator program. R.E.M.G. is grateful for continuing Guest Worker privileges at the DAO (National Research Council of Canada) and for observing time on the DAO 1.2 m telescope. REFERENCES Adelman, S. J., Caliskan, H., Gulliver, A. F., & Teker, A. 2006, A&A, 447, 685 Bidelman, W. P. 1981, AJ, 86, 553 Bidelman, W. P. 1985, AJ, 90, 341 Böhm-Vitense, E., Carpenter, K., Robinson, R., Ake, T., & Brown, J. 2000, ApJ, 533, 969 Bond, H. E. 1984, in NASA Conf. Publ. 2349, Future of Ultraviolet Astronomy Based on Six Years of IUE Research, ed. J. M. Mead, R. D. Chapman, & Y. Kondo (Washington, DC: NASA), 289 Bond, H. E., Pollacco, D. L., & Webbink, R. F. 2003, AJ, 125, 260 Cardelli, J. A., Clayton, G. C., & Mathis, J. S. 1989, ApJ, 345, 245 Castelli, F., & Kurucz, R. L. 2003, in IAU Symp. 210, Modeling of Stellar Atmospheres, ed. N. E. Piskunov, W. W. Weiss, & D. F. Gray (San Francisco, CA: ASP), Poster A20 on Enclosed CD-ROM Catalán, S., Isern, J., Garcia-Berro, E., & Ribas, I. 2009, J. Phys. Conf. Ser., 172, 012007 Crawford, D. L., & Mandwewala, N. 1976, PASP, 88, 917 Dominguez, I., Chieffi, A., Limongi, M., & Straniero, O. 1999, ApJ, 524, 226 Edvardsson, B., Andersen, J., Gustafsson, B., Lambert, D. L., Nissen, P. E., & Tomkin, J. 1993, A&A, 275, 101 Gray, R. O., & Corbally, C. J. 1994, AJ, 107, 742 Gray, R. O., Corbally, C. J., Garrison, R. F., McFadden, M. T., Bubar, E. J., McGahee, C. E., O’Donoghue, A. A., & Knox, E. R. 2006, AJ, 132, 161 Gray, R. O., Corbally, C. J., Garrison, R. F., McFadden, M. T., & Robinson, P. E. 2003, AJ, 126, 2048 Gray, R. O., Graham, P. W., & Hoyt, S. R. 2001, AJ, 121, 2159 Gray, R. O., & Griffin, R. E. M. 2007, AJ, 134, 96 Gulliver, A. F., Hill, G., & Adelman, S. J. 1996, in ASP Conf. Ser. 108, M.A.S.S. Model Atmospheres and Spectrum Synthesis, ed. S. J. Adelman, F. Kupka, & W. W. Weiss (San Francisco, CA: ASP), 232 Holberg, J. B., Barstow, M. A., & Burleigh, M. R. 2003, ApJS, 147, 145 Kurucz, R. 1993, ATLAS9 Stellar Atmosphere Programs and 2 km/s Grid, Kurucz CD-ROM No. 13 (Cambridge, MA: Smithsonian Astrophysical Observatory), 13 Lejeune, T., & Schaerer, D. 2001, A&A, 366, 538 Makarov, V. V., & Kaplan, G. H. 2005, AJ, 129, 2420 Mamajek, E. E., & Hillenbrand, L. A. 2008, ApJ, 687, 1264 Martin, D. C., et al. 2005, ApJ, 619, L1 McClure, R. D. 1983, ApJ, 268, 264 McClure, R. D., Fletcher, J. M., & Nemec, J. M. 1980, ApJ, 238, 35 Morrissey, P., et al. 2007, ApJS, 173, 682 Moultaka, J., Ilovaisky, S. A., Prugniel, P., & Soubiran, C. 2004, PASP, 116, 693 North, P., Berthet, S., & Lanz, T. 1994, A&A, 281, 775 North, P., & Duquennoy, A. 1991, A&A, 244, 335 North, P., Jorissen, A., & Mayor, M. 2000, in IAU Symp. 177, The Carbon Star Phenomenon, ed. R. F. Wing (Dordrecht: Kluwer), 269 North, P., & Lanz, T. 1991, A&A, 251, 489 Parsons, S. B. 1983, ApJS, 53, 553 Rufener, F., & Maeder, A. 1971, A&AS, 4, 43 Schlegel, D. J., Finkbeiner, D. P., & Davis, M. 1998, ApJ, 500, 525 Soderblom, D. R., Duncan, D. K., & Johnson, D. R. H. 1991, ApJ, 375, 722 Tomkin, J., Lambert, D. L., Edvardsson, B., Gustafsson, B., & Nissen, P. E. 1989, A&A, 219, L15 van Leeuwen, F. 2007, A&A, 474, 653 Vauclair, S., & Vauclair, G. 1982, ARA&A, 20, 37 Vaughan, A. H., Preston, G. W., & Wilson, O. C. 1978, PASP, 90, 267

6. CONCLUSIONS In this paper, we have analyzed both guest investigator (GI) and archived GALEX FUV observations of Ba dwarfs and have demonstrated for the first time that the Ba dwarfs as a group show a significant FUV excess when compared to a control sample of F-type dwarfs. That FUV excess can be interpreted as the first direct detection of WD companions of Ba dwarfs. The authors acknowledge the useful comments of an anonymous referee which helped to significantly improve this paper.

8