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1996MNRAS.278..688S Mon. Not. R. Astron. Soc. 278, 688-696 (1996) The chemical composition of IK Pegasi B. Smalley/* K. C. Smith,2 D. Wonnacote and C. S. Allen2 lDepartment of Physics, Keele University, Keele, Staffordshire, ST5 5BG Department of Physics and Astronomy, University College London, Gower Street, London, WCIE 6BT 3Mullard Space Science Laboratory, University College London, Holmbury St. Mary, Dorking, Surrey, RH5 6NT 2 Accepted 1995 August 30. Received 1995 August 23; in original form 1995 March 30 ABSTRACT A detailed abundance analysis of the pulsating A-type star IK Peg A is presented. It is found that the Ca and Sc abundances are approximately solar and the Fe-group elements slightly enhanced. Hence, IK Peg A is not a classical Am star, but the results are not inconsistent with its spectroscopic classification as a marginal Am star. The presence of a massive white dwarf companion (IK Peg B) indicates that the binary system has previously undergone a common envelope phase. Enhancements of Ba and Sr are found, which may be evidence of mass transfer from the white dwarf progenitor during this common envelope phase. This, however, is purely speculative, since the abundance anomalies may be explained by radiative diffusion processes operating in the atmosphere of IK Peg A, even though it is undergoing smallamplitude pulsations. It is suggested that IK Peg A is a hot member of the F str M077 stars. Key words: stars: abundances - stars: chemically peculiar - stars: individual: IK Peg A - stars: rotation. 1 INTRODUCTION The sixth magnitude A-type star IK Pegasi (HR 8210, HD 204188) has been the subject of much interest over the last few years. It is a well-known single-lined spectroscopic binary with a period of 21.7 d (Harper 1935; Batten, Fletcher & MacCarthy 1989). Only recently has the companion (IK Peg B) been positively identified as a massive ( ~ 1.15 Mo) hot (Tefl ~ 35 000 K) white dwarf (Wonnacott, Kellett & Stickland 1993; Landsman, Simon & Bergeron 1993). While much of this interest has concentrated on the white dwarf (Barstow et al. 1994; Barstow, Holberg & Koester 1994; Landsman, Simon & Bergeron 1995), the primary itself (IK Peg A) is by no means an uninteresting main-sequence A-type star. It is known to be undergoing small-amplitude pulsations (Kurtz 1978; Wonnacott et al. 1994) and is currently believed to exhibit mild spectroscopic pecularities similar to those associated with the metalliclined A-type stars (Cowley et al. 1969; Abt & Bidelman 1969). The metallic-lined (Am) stars are a spectroscopic class of A-type stars in which the spectral type inferred from the metal lines is at least five spectral subtypes later than that * Formerly the Department of Physics and Astronomy, University College London, Gower Street, London, WC1E 6BT. inferred from the calcium Hand K lines (Titus & Morgan 1940; Roman, Morgan & Eggen 1948). The hydrogen-line spectral type is intermediate between the two. This defines the classical Am stars. A marginal Am star is one in which the metallic and calcium types differ by less than five subtypes. These definitions only describe the appearance of the spectrum and do not imply anything about the abundances of the elements. However, based on a review of the contemporary detailed abundance analyses of Am stars, Conti (1970) proposed a new definition of the Am phenomenon: 'The Am phenomenon is present in stars that have an apparent surface underabundance of Ca (and/or Sc) and/or apparent overabundance of the Fe group and heavier elements'. Within this definition, a classical Am star has underabundances of Ca and Sc and an overabundance of Fe-group elements. Subsequent studies of Am stars have shown that the abundance anomalies are the results of radiative diffusion in the stable atmospheres of these slowly rotating stars (Michaud 1970). Like normal A-type stars, the Tefl of Am stars can be reliably determined from the Balmer lines (Smalley & Dworetsky 1993). The spectroscopic classification of IK Peg A is uncertain. It is most often cited as a marginal Am star (Cowley et al. 1969; Kurtz 1978), but was called 'definitely Am' by Abt & Bidelman (1969). This uncertainty was discussed by Wonnacott et al. (1994), who used medium-resolution spectra to © 1996 RAS © Royal Astronomical Society • Provided by the NASA Astrophysics Data System 1996MNRAS.278..688S The chemical composition of IK Pegasi 689 show that the mean metal abundance does not differ significantly from that of the Sun. Any anomalies are therefore subtle and are likely to be detected only at high-resolution. There have been relatively few abundance studies of IK Peg A. Cowley & Aikman (1980) extended the method of wavelength coincidence statistics (Cowley & Henry 1979) in an attempt to estimate abundances using empirical calibrations. They found that Mn, Y and Fe were a few tenths of a dex underabundant and the Cr was strongly depleted. However, this method is mainly intended for use in the preliminary estimation of abundances and is not a substitute for detailed abundance analysis based on equivalent-width measurements or spectrum synthesis. Guthrie (1987), in his study of the calcium abundances in Am stars, found that Ca was marginally overabundant in IK Peg. This was discussed by Wonnacott et a1. (1994), who concluded that the Ca abundance is solar to within ± 0.2 dex. In this paper, a detailed element-by-element abundance analysis is presented in order to determine the atmospheric abundance pattern of IK Peg A. 2 OBSERVATIONS AND REDUCTIONS During two observing runs with the Hamilton Echelle Spectrograph on the Coude Auxiliary Telescope at Lick Observatory, California, blue- and red-region echelle spectra were obtained. The spectrograph has a resolving power of 48 000 and a mean linear dispersion of 2.54 A mm- 1 (Vogt 1987). During the first run, two blue-region frames were obtained, giving complete coverage from 3900 to 4800 A. A Texas Instruments 800 x 800 pixel thinned, back-illuminated CCD was used. During the second run, four red-region frames were obtained on consecutive mornings just prior to the beginning of twilight. A Ford Aerospace 2048 x 2048 pixel unthinned, front-illuminated, phosphor-coated CCD chip was used, giving complete wavelength coverage from 4000 to 9000 A. The four frames were taken in order to monitor the night-to-night changes in radial velocity, as well as to increase the signal-to-noise ratio in the final co-added spectrum. The majority of the data reduction was performed using the Lick VISTA and Starlink FIGARO packages (Stover 1988; Meyerdierks 1993). The echelle images suffer from general scattered light, which was successfully removed using the procedures outlined in the VISTA Cookbook (Pogge, Goodrich & Veilleux 1988) and discussed by Allen (in preparation). That this process is adequate was demonstrated by performing the same reduction procedures on a solar spectrum. The instrumental profile was measured from the Th-Ar comparison arc spectra and convolved with the standard solar flux atlas from Kitt Peak (Kurucz et a1. 1984). A comparison between this spectrum and the data yielded no significant differences. In addition, the instrumental profile was found to be very slightly asymmetric. Fortunately, the rotational velocity of IK Peg A was sufficiently high to render the effects of this asymmetry negligible (Section 3.3). The overall signal-to-noise ratio (SIN) of the blue-region spectra was around 100:1. The red region, however, was slightly more noisy and the SIN varied with wavelength, due to variations in stellar flux levels and in the instrumental and CCD responses. The majority of the red-region spectra had an SIN of at least 70:1. 3 ABUNDANCE AN ALYSIS The basic atmospheric parameters T eff, log g and [M/H] are prerequisities for a detailed abundance analysis. These parameters were fully discussed by Wonnacott et a1. (1994) and are adopted here (Table 1). Consequently, a Kurucz (1979) solar-composition model atmosphere with Teff = 7770 K and log g = 4.25 was used in the analysis. The analysis was performed using the LTE spectrum synthesis code UCLSYN developed by Smith (1992). The first step in the analysis is the determination of microturbulence (~t) and the projected rotation velocity (v sin i). These are discussed in Sections 3.2 and 3.3, respectively. 3.1 Equivalent-width measurements The Kurucz & Pevtremann (1975) and Kurucz (1988, private communication) line lists were used to identify absorption lines which appeared to be unblended (i.e. more than 95 per cent in the absorption feature was due to only one line). This process was complicated by the moderately high value of v sin i which caused many otherwise single lines to be blended together. Nevertheless, a few hundred lines were identified and their equivalent widths (W,J measured (see Table 2). In addition, several blended lines were identified. These were required to obtain abundances of interesting elements or to increase the number of lines of important elements. In these cases, the equivalent width of the blend was measured and the abundances determined using spectrum synthesis (see Section 4). A literature search was made to find a more accurate source of log gfvalues for the measured lines and blending components. For the iron-group elements, the damping constants were taken from Kurucz (1988, private communication); for all other elements Kurucz' WIDTH defaults were adopted. 3.2 Microturbulence Microturbulence (~t) is a fitting parameter, orginally introduced to make abundance results from weak lines agree with those from strong lines. Its value is crucial to the accurate determination of elemental abundances, since if too low a value is used, the abundances will be overestimated and vice versa. In spite of the importance of microturbuTable 1. Basic atmospheric parameters of IK Peg A. Teff logg [M/H] ~t vsini 7770 ± 100 K 4.25 ± 0.10 +0.07 ± 0.20 2.6 ± 0.2 km 8- 1 32.5 ± 2.5 km S-l Notes: Tefl, log g and [M/H] are taken from Wonnacott et al. (1994); ~t and v sin i are determined in the present. © 1996 RAS, MNRAS 278, 688-696 © Royal Astronomical Society • Provided by the NASA Astrophysics Data System 1996MNRAS.278..688S 690 B. Smalley et al. Table 2. Individual line identifications, equivalents widths and abundances. ID CI CI CI CI CI CI CI CI CI CI CI CI CI CI CI CI CI CI NI NI NI 01 01 01 Nal Nal MgI MgI MgI All All AlI All Sil Sil Sil Sil Sil Sil Sil Sil Sil Sil Sil Sin Si n Sin Si n Sin SI SI KI Cal Cal Cal Cal Cal Cal Cal Cal Cal Cal Cal Cal Cal Cal Cal Cal Cal Cal Cal Cal A X loggf 4268.99 4371.33 4762.31 4762.53 4766.67 4770.03 4771.74 4775.90 4817.33 4932.00 5052.17 5380.34 6587.61 7087.83 7111.48 7113.18 7115.17 8335.15 7468.31 8216.34 8223.13 6158.19 7947.55 7950.80 4664.81 4978.54 4702.99 5528.40 8213.01 3944.01 3961.52 6696.02 6698.67 5772.15 5793.07 5797.86 5948.54 7071.82 7405.77 7680.27 7742.72 7918.39 7932.35 7944.00 4130.88 5055.98 5056.31 6347.11 6371.37 6748.79 6757.16 7698.97 4318.65 4425.44 4434.96 4435.69 4578.56 4581.40 4585.86 4685.27 5581.97 5588.75 5594.46 5857.45 6122.22 6161.29 6162.17 6439.07 6449.81 6462.57 6717.68 7148.15 7.68 7.68 7.48 7.48 7.48 7.48 7.49 7.49 7.48 7.68 7.68 7.68 8.54 8.65 8.64 8.65 8.64 7.68 10.34 10.34 10.33 10.74 12.54 12.54 2.10 2.10 4.35 4.35 5.75 0.00 0.01 3.15 3.14 5.08 4.93 4.95 5.08 5.95 5.61 5.86 6.21 5.95 5.96 5.98 9.84 10.07 10.07 8.12 8.12 7.87 7.87 0.00 1.90 1.88 1.89 1.89 2.52 2.52 2.53 2.93 2.52 2.53 2.52 2.93 1.89 2.52 1.90 2.53 2.52 2.52 2.71 2.71 -2.360 -2.080 -2.280 -2.200 -2.400 -2.280 -1.700 -2.270 -2.530 -1.780 -1.240 -1.570 -1.050 -1.480 -1.070 -0.760 -1.030 -0.460 -0.270 0.160 -0.240 -0.332 0.500 0.340 -1.550 -1.206 -0.380 -0.480 -0.530 -0.644 -0.345 -1.343 -1.650 -1.380 -1.480 -1.190 -1.240 -1.330 -0.570 -0.560 -0.690 -0.590 -0.450 -0.380 0.483 0.517 -0.437 0.225 -0.074 -0.530 -0.240 -0.168 -0.208 -0.385 -0.029 -0.500 -0.560 -0.337 -0.186 -0.880 -0.710 0.210 -0.050 0.230 -0.409 -1.020 -0.218 0.470 -0.550 0.310 -0.610 0.208 W", 15 20 20 23 12 16 67 30 9 38 66 51 33 32 30 34 33 148 36 48 22 29 22 17 10 20 159 160 86 175 180 13 10 27 32 38 64 20 65 36 46 55 72 53 74 62 13 155 118 31 40 48 100 85 96 71 26 31 45 8 34 110 116 97 108 26 119 124 54 119 53 112 (ELjH) 8.41 ± 0.25 8.30 ± 0.25 8.32 ± 0.25 8.31 ± 0.24 8.19 ± 0.35 8.21 ± 0.25 8.46 ± 0.23 8.53 ± 0.14 8.21 ± 0.31 8.31 ± 0.23 8.12 ±' 0.14 8.27 ± 0.16 8.20 ± 0.15 8.58 ± 0.17 8.13 ± 0.18 7.89 ± 0.17 8.14±0.17 8.26 ± 0.16 8.15 ± 0.18 7.95 ± 0.26 7.85 ± 0.28 8.63 ± 0.14 8.92 ± 0.22 8.94 ± 0.25 6.41 ± 0.21 6.39 ± 0.16 7.30 ± 0.13 7.74 ± 0.25 7.71 ± 0.24 6.30 ± 0.14 6.18 ± 0.13 6.44 ± 0.22 6.62 ± 0.27 7.46 ± 0.25 7.54 ± 0.25 7.37 ± 0.24 7.88 ± 0.24 7.86 ± 0.29 7.59 ± 0.16 7.34 ±0.26 7.84 ± 0.26 7.71 ± 0.25 7.80 ± 0.25 7.50 ± 0.25 7.76 ± 0.14 7.67 ± 0.16 7.61 ± 0.23 7.87 ± 0.15 7.71 ± 0.15 7.27 ± 0.15 7.13 ± 0.15 4.88 ± 0.14 6.42 ± 0.16 6.26 ± 0.15 6.08 ± 0.15 6.18 ± 0.14 5.99 ± 0.15 5.86 ± 0.15 5.95 ± 0.14 6.00 ± 0.30 6.24 ± 0.25 6.44 ± 0.25 6.80 ± 0.26 6.49 ± 0.25 6.50 ± 0.17 6.38 ± 0.25 6.48 ± 0.17 6.37 ± 0.26 6.37 ± 0.24 6.45 ± 0.26 6.54 ± 0.24 6.56 ± 0.27 Ref [21) [21) [21) [21) [21) [21) [21) [5) [21) [20) [5) [5) [5) [5) [5) [5) [15) [5) [20) [11) [11) [22) l20) [20) [22) [22) [20) [22) [11) [14) [14) [22) [22) [22) [22) [22) [22) [11) [18) [22) {11) [22) [22) [22) [22) [1) [1) [22) [22) [6) [6] [22) [22) [22) [22) [22) [22) [22) [22) [22) [22) [22) [22) [22) [22) [22) [22) [22] [22) [22) [22) [12] ID Sen Sen Sen Sen Se n Sen Sen Til Til Til Tin Till Tin Tin Tin Till Tin Tin Tin Tin Tin Tin Tin Tin Tin Tin VI Vn Vn VII Vn Vn Vn Vn Crl Cr! Crl Crl Cr! Cr! Crl Cr! Crl Crn Crll Crn Crn Crn Crn Crn Crn Crn Crn Mnl Mnl Mnl Mnl Mnl Mnl Mnl Mnl Mnl Fel Fel Fel Fel Fel Fel Fel Fel Fel Fel A 4320.74 4420.66 4670.40 5031.02 5239.81 5526.81 6604.60 4552.45 4991.07 4999.50 4012.39 4028.34 4053.83 4312.86 4386.86 4464.46 4501.27 4563.76 4589.96 4779.99 4805.09 5072.28 5129.15 5154.07 5185.91 6827.91 4379.23 3916.41 4023.39 4035.63 4036.78 4183.44 4202.36 4528.50 4254.33 4496.86 4646.17 4718.43 4922.27 4954.81 5204.51 5206.04 5791.01 4242.36 4261.91 4558.66 4588.22 4634.07 4848.23 4876.40 5237.33 5313.56 5334.87 4030.75 4034.48 4083.63 4451.59 4453.01 4754.04 4783.42 6013.48 6021.79 3922.91 4021.87 4071.74 4073.76 4147.67 4200.92 4202.03 4213.65 4217.55 4222.22 X loggf 0.60 0.62 1.36 1.36 1.46 1.77 1.36 0.84 0.84 0.83 0.57 1.89 1.89 1.18 2.60 1.16 1.12 1.22 1.24 2.05 2.06 3.12 1.89 1.57 1.89 3.10 0.30 1.43 1.80 1.79 1.48 2.05 1.70 2.28 0.00 0.94 1.03 3.19 3.10 3.12 0.94 0.94 3.32 3.87 3.86 4.07 4.07 4.07 3.86 3.85 4.07 4.07 4.07 0.00 0.00 2.16 2.89 2.94 2.28 2.30 3.07 3.08 0.05 2.76 1.61 3.27 1.49 3.40 1.49 2.85 3.43 2.45 -0.260 -2.140 -0.370 -0.260 -0.770 0.130 -1.480 -0.340 0.380 0.250 -1.610 -1.000 -1.210 -1.160 -1.260 -2.080 -0.750 -0.960 -1.790 -1.370 -1.100 -0.750 -1.390 -1.920 -1.350 -1.579 0.580 -1.060 -0.518 -0.622 -1.540 -0.946 -1.750 -1.098 -0.114 -1.150 -0.700 0.090 0.270 -0.300 -0.208 0.019 0.324 -1.160 -1.360 -0.460 -0.630 -1.020 -1.150 -1.450 -1.160 -1.650 -1.562 -0.470 '-0.811 -0.250 0.278 -0.490 -0.086 0.042 -0.251 0.034 -1.651 -0.660 -0.022 -0.920 -2.104 -1.000 -0.708 -1.290 -0.510 -0.967 W", 88 10 41 60 24 62 13 30 52 44 133 115 66 135 74 87 167 172 103 83 113 75 90 73 72 10 20 59 67 61 16 29 24 22 148 38 57 18 31 29 98 143 32 90 92 146 122 91 98 66 99 64 59 125 120 26 66 13 52 69 31 31 146 114 181 65 84 36 153 67 103 103 (ELjH) 2.82 ± 0.24 3.25 ± 0.28 2.75 ± 0.23 2.88 ± 0.24 2.88 ± 0.26 2.79 ± 0.24 3.17 ± 0.29 5.47 ± 0.25 5.08 ± 0.22 5.08 ± 0.16 5.39 ± 0.30 5.40 ± 0.26 4.81 ± 0.24 5.36 ± 0.28 5.45 ± 0.24 5.36 ± 0.24 5.46 ± 0.32 5.82 ± 0.32 5.36 ± 0.25 5.23 ± 0.24 5.44 ± 0.26 5.28 ± 0.24 5.20 ± 0.26 5.25 ± 0.26 4.91 ± 0.24 4.80 ± 0.41 3.95 ± 0.17 4.16 ± 0.10 4.00 ± 0.23 4.00 ± 0.23 3.87 ± 0.14 3.99 ± 0.24 4.43 ± 0.24 4.13 ± 0.15 5.95 ± 0.22 5.72 ± 0.12 5.61 ± 0.11 5.70 ± 0,25 5.74 ± 0.24 6.28 ± 0.24 5.60 ± 0.15 6.16 ± 0.20 5.82 ± 0.25 5.73 ± 0.16 5.96 ± 0.15 6.13 ± 0.20 5.85 ± 0.18 5.70 ± 0.15 5.97 ± 0.16 5.71 ± 0.14 5.91 ± 0.25 5.89 ± 0.24 5.74 ± 0.24 5.38 ± 0.22 5.60 ± 0.20 5.07 ± 0.16 5.58 ± 0.10 5.44 ± 0.16 5.33 ± 0.11 5.45 ± 0.11 5.68 ± 0.17 5.39 ± 0.17 7.57 ± 0.23 7.82 ± 0.18 7.64 ± 0.19 7.58 ± 0.15 7.72 ± 0.13 7.27 ± 0.15 7.67 ± 0.20 7.63 ± 0.15 7.85 ± 0.17 7.60 ± 0.15 Ref [13) [19) [13) [13) [13) [13) [13) [13) [13) [13) [13) [13) [13) [13) [13) [13) [13) [13) [13) [13) [13) [13) [13) [13) [13) [12) [13) [13) [12) [12) [13) [12) [13) [4) [13) [13) [13J [13) [13) [13] [13) [13) [12) [17] [17) [17) [17) [17) [17) [17) [13) [13) [12) [13) [13) [13) [13] [13] [13) [13) [13) [13) [7) [7) [7) [7) [7) [7] [7] [7) [7) [7) © 1996 RAS, MNRAS 278, 688-696 © Royal Astronomical Society • Provided by the NASA Astrophysics Data System 1996MNRAS.278..688S The chemical composition of IK Pegasi Table 2 - continued ID Fe I Fe I Fel Fe I Fe I Fe I Fe I Fel Fel Fel Fel Fe I Fel Fel Fel Fe I Fe I Fel Fe I Fe I Fel Fe I Fel Fel Fel Fe I Fel Fel Fel Fe I Fe I Fe I Fel Fel Fel Fe I Fe I Fe I Fel Fel Fel Fe I Fe I Fe I Fe I Fel Fe I Fe I Fel Fel Fe I Fe I Fe I Fel Fel Fel Fe I Fe I Fel Fe I Fe I Fe I Fel Fel Fel Fel Fel Fe I Fel Fe I Fe I Fe I A 4285.44 4466.55 4476.02 4484.23 4485.68 4494.57 4531.15 4607.65 4611.28 4643.47 4647.44 4691.41 4728.55 4745.80 4768.32 4903.32 4920.51 4930.33 4942.46 4946.39 4966.10 4973.11 4988.96 5001.87 5022.24 5065.01 5068.77 5074.76 5090.79 5097.00 5123.72 5133.69 5137.39 5162.29 5232.95 5281.80 5302.31 5324.19 5353.39 5367.47 5369.97 5389.46 5393.17 5400.51 5434.53 5445.04 5466.40 5480.87 5497.52 5506.78 5565.70 5569.62 5572.85 5586.76 5633.97 5662.52 5679.02 5705.99 5717.85 5816.36 5862.35 5905.67 5934.66 5984.81 5987.07 6003.03 6020.17 6024.07 6055.99 6065.48 6078.49 6230.73 X 3.24 2.83 2.85 3.60 3.69 2.20 1.49 3.27 3.65 3.65 2.95 2.99 3.65 3.65 3.69 2.88 2.83 3.96 4.22 3.37 3.33 3.96 4.15 3.88 3.98 4.26 2.94 4.22 4.26 4.28 1.01 4.18 4.18 4.18 2.94 3.04 3.28 3.21 4.10 4.42 4.37 4.42 3.24 4.37 1.01 4.39 4.37 4.22 1.01 0.99 4.61 3.42 3.40 3.37 4.99 4.18 4.65 4.61 4.28 4.55 4.55 4.65 3.93 4.73 4.79 3.88 4.61 4.55 4.73 2.61 4.79 2.56 loggf -1.190 -0.590 -0.726 -0.720 -1.020 -1.136 -2.155 -1.545 -0.699 -1.290 -1.310 -1.450 -1.442 -0.790 -1.109 -1.080 0.060 -1.350 -1.243 -1.170 -0.890 -0.950 -0.890 0.010 -0.530 -0.134 -1.230 -0.200 -0.400 -0.277 -3.068 0.140 -0.400 0.020 -0.190 -1.020 -0.880 -0.240 -0.840 0.350 0.350 -0.410 -0.910 -0.160 -2.122 -0.020 -0.630 -1.260 -2.849 -2.797 -0.285 -0.540 -0.310 -0.210 -0.270 -0.541 -0.920 -0.530 -1.130 -0.680 -0.058 -0.730 -1.170 -0.343 -0.556 -1.120 -0.270 -0.120 -0.460 -1.530 -0.424 -1.281 WA 46 131 98 71 43 96 73 44 71 34 42 53 23 44 39 83 146 26 27 67 82 49 52 109 74 100 87 98 45 74 55 116 84 121 148 98 93 135 48 114 127 48 90 79 101 96 61 25 69 71 62 95 125 132 46 74 20 42 31 36 66 37 36 48 30 51 76 95 45 66 39 100 {EL/H} 7.51 ± 0.15 8.02 ± 0.20 7.54 ± 0.26 7.65 ± 0.24 7.59 ± 0.14 7.41 ± 0.13 7.54 ± 0.12 7.82 ± 0.24 7.65 ± 0.14 7.68 ± 0.15 7.32 ± 0.24 7.65 ± 0.24 7.60 ± 0.25 7.34 ± 0.24 7.60 ± 0.24 7.64 ± 0.15 7.51 ± 0.19 7.79 ± 0.15 7.89 ± 0.24 7.83 ± 0.24 7.74 ± 0.15 7.78 ± 0.14 7.91 ± 0.14 7.67 ± 0.18 7.74 ± 0.16 7.91 ± 0.25 7.86 ± 0.16 7.87 ± 0.15 7.38 ± 0.15 7.69 ± 0.25 7.80 ± 0.13 7.73 ± 0.25 7.86 ± 0.16 7.80 ± 0.24 7.81 ± 0.20 7.88 ± 0.17 7.84 ± 0.17 7.84 ± 0.19 7.74 ± 0.24 7.76 ± 0.17 7.92 ± 0.18 7.54 ± 0.24 7.79 ± 0.16 7.70 ± 0.24 7.49 ± 0.15 7.76 ± 0.24 7.92 ± 0.24 7.84 ± 0.17 7.75 ± 0.13 7.71 ± 0.13 7.76 ± 0.24 7.61 ± 0.17 7.85 ± 0.19 7.84 ± 0.18 7.78 ± 0.15 7.86 ± 0.25 7.69 ± 0.19 7.70 ± 0.16 7.87 ± 0.16 7.70 ± 0.25 7.51 ± 0.24 7.84 ± 0.16 7.74 ± 0.16 7.69 ± 0.24 7.65 ± 0.25 7.89 ± 0.24 7.92 ± 0.24 8.03 ± 0.26 7.76 ± 0.24 7.56 ± 0.12 7.67 ± 0.24 7.76 ± 0.14 Ref [7] [7] [12] [7] [7] [7] [7] [12] [10] [7] [7] [7] [12] [12] [12] [7] [7] [7] [12] [7] [7] [7] [7] [7] [7] [12] [7] [7] [7] [12] [7] [7] [7] [7] [7] [7] [7] [7] [7] [7] [7] [7] [7] [7] [7] [7] [7] [7] [7] [7] [12] [7] [7] [7] [7] [12] [7] [7] [7] [7] [12] [7] [7] [12] [12] [7] [7] [7] [7] [7] [12] [7] ID Fel Fel Fel Fel Fel Fel Fel Fel Fel Fel Fel Fel Fel Fel Fel Fel Fel Fel Fel Fel Fel Fel Fell Fell Fell Fell Fell Fell Fell Fell Fell Fell Fell Fell Fell Fell Fell COl COl COl Nil Nil Nil Nil Nil Nil Nil Nil Nil Nil Nil Nil Nil Nil Nil Nil Nil Nil CUI CUI Znl Znl Srll Srll Sr II Sr II YII YII YII YII YII YII A 6252.56 6256.37 6265.14 6393.60 6400.01 6411.66 6430.85 6633.75 6677.99 6843.65 6855.16 7068.40 7090.38 7130.92 7445.75 7495.06 7511.02 7832.19 7937.13 8046.05 8085.18 8327.05 4491.40 4508.28 4515.34 4520.23 4541.52 4576.33 4620.51 4629.34 4663.70 4731.44 4923.92 5254.93 5362.87 5991.38 6432.68 4092.39 4118.77 4867.87 4686.22 4714.42 4715.78 4752.43 4807.00 4829.03 4831.18 4904.41 4980.16 5115.39 5146.48 5715.07 6643.63 7122.19 7522.76 7525.11 7555.60 7797.59 5105.54 5782.13 4722.16 4810.53 4077.71 4161.79 4215.52 4305.44 3950.36 4358.73 4398.01 4883.69 4900.11 5087.42 X 2.40 2.45 2.18 2.43 3.60 3.65 2.18 4.56 2.69 4.55 4.56 4.08 4.23 4.22 4.26 4.22 4.18 4.43 4.31 4.42 4.45 2.20 2.85 2.85 2.84 2.81 2.85 2.84 2.83 2.81 2.89 2.89 2.89 3.23 3.20 3.15 2.89 0.92 1.05 3.12 3.60 3.38 3.54 3.66 3.68 3.54 3.61 3.54 3.61 3.83 3.71 4.09 1.68 3.54 3.66 3.63 3.85 3.90 1.38 1.64 4.01 4.06 0.00 2.94 0.00 3.04 0.10 0.10 0.13 1.08 1.03 1.08 loggf -1.687 -2.620 -2.550 -1.620 -0.520 -0.820 -2.006 -0.780 -1.470 -0.930 -0.485 -1.380 -1.210 -0.790 -0.237 -0.102 0.107 0,018 0.152 -0.082 -0.240 -1.525 -2.700 -2.210 -2.480 -2.600 -3.050 -3.040 -3.280 -2.370 -3.820 -3.360 -1.320 -3.227 -2.739 -3.740 -3.740 -0.940 -0.490 0.226 -0.640 0.230 -0.340 -0.700 -0.640 -0.330 -0.420 -0.170 -0.110 -0.110 -0.060 -0.352 -2.200 0.040 -0.575 -0.546 -0.046 -0.262 -1.510 -1.782 -0.390 -0.170 0.210 -0.470 -0.180 -0.110 -0.490 -1.360 -1.000 0.070 -0.090 -0.170 W>. 71 28 20 74 109 83 52 40 76 31 38 22 24 48 65 82 119 90 103 90 91 113 124 152 135 130 105 114 81 143 39 96 210 73 129 65 67 20 49 9 22 107 37 41 29 59 39 62 56 52 77 26 20 66 53 53 71 36 14 7 37 60 229 46 220 72 59 24 49 53 53 41 {EL/H} 7.63 ± 0.12 7.94 ± 0.26 7.49 ± 0.17 7.62 ± 0.25 7.91 ± 0.26 7.86 ± 0.25 7.51 ± 0.13 7.87 ± 0.16 7.69 ± 0.25 7.84 ± 0.17 7.53 ± 0.25 7.75 ± 0.21 7.74 ± 0.20 7.74 ± 0.18 7.45 ± 0.26 7.53 ± 0.26 7.82 ± 0.27 7.68 ± 0.26 7.64 ± 0.27 7.76 ± 0.26 7.96 ± 0.26 7.86 ± 0.18 7.80 ± 0.19 7.79 ± 0.30 7.76 ± 0.29 7.76 ± 0.27 7.79 ± 0.26 7.92 ± 0.26 7.61 ± 0.24 7.75 ± 0.29 7.59 ± 0.23 7.95 ± 0.24 7.66 ± 0.20 7.70 ± 0.24 8.03 ± 0.28 8.01 ± 0.24 7.82 ± 0.24 5.05 ± 0.19 5.18 ± 0.16 5.20 ± 0.30 6.35 ± 0.25 6.71 ± 0.26 6.30 ± 0.24 6.81 ± 0.24 6.56 ± 0.24 6.65 ± 0.24 6.49 ± 0.24 6.52 ± 0.24 6.40 ± 0.24 6.50 ± 0.24 6.71 ± 0.24 6.46 ± 0.26 6.37 ± 0.27 6.26 ± 0.26 6.79 ± 0.26 6.73 ± 0.26 6.65 ± 0.26 6.38 ± 0.27 4.23 ± 0.22 4.33 ± 0.37 4.59 ± 0.15 4.76 ± 0.14 3.35 ± 0.19 3.24 ± 0.23 3.60 ± 0.19 3.35 ± 0.15 2.19 ± 0.14 2.43 ± 0.16 2.50 ± 0.14 2.21 ± 0.14 2.30 ± 0.14 2.21 ± 0.15 Ref [7] [7] [7] [7] [7] [7] [7] [7] [7] [7] [12] [7] [7] [7] [12] [12] [12] [12] [12] [12] [12] [7] [7] [7] [7] [7] [7] [7] [7] [7] [7] [7] [7] [12] [12] [7] [7] [7] [7] [12] [7] [7] [7] [7] [7] [7] [7] [7] [7] [7] [12] [12] [7] [7] [12] [12] [12] [12] [20] [20] [2] [2] [16] [20] [16] [16] [9] [8] [9] [9] [9] [9] © 1996 RAS, MNRAS 278, 688-696 © Royal Astronomical Society • Provided by the NASA Astrophysics Data System 691 1996MNRAS.278..688S 692 B. Smalley et al. Table 2 - continued ID Yn Yn YII ZrIl ZrIl ZrIl ZrlI ZrIl Zru A 5123.21 5200.41 7881.88 3991.13 4048.67 4149.20 4208.98 4211.90 4359.74 X 0.99 0.99 1.84 0.75 0.80 0.80 0.71 0.53 1.24 loggf W). -0.830 -0.570 -0.570 -0.300 -0.480 -0.030 -0.460 -0.980 -0.560 26 24 8 35 31 41 26 9 16 (EL/H) 2.53 ± 0.17 2.22 ± 0.17 2.22 ± 0.37 2.56 ± 0.14 2.69 ± 0.14 2.40 ± 0.14 2.48 ± 0.15 2.33 ± 0.25 2.74 ± 0.18 Ref [9] [9] [9] [8] [8] [3] [3] [8] [8] ID ZrIl BaIl BaIl Ball BaIl Ball Ball Ndu Ndu A 4496.97 4524.93 4554.03 4899.93 5853.67 6141. 72 6496.90 4012.30 4061.10 X 0.71 2.51 0.00 2.72 0.60 0.70 0.60 0.63 0.47 loggf W>. -0.810 -0.350 0.163 -0.170 -1.010 -0.077 -0.377 0.420 0.290 7 44 195 50 92 167 153 20 32 (EL/H) 2.15 ± 0.28 3.20 ± 0.24 3.41 ± 0.14 3.26 ± 0.10 3.02 ± 0.14 3.44 ± 0.21 3.39 ± 0.25 1.99 ± 0.17 2.26 ± 0.16 Ref [3] [20] [20] [20] [20] [20] [20] [20] [20] Notes: .Ie is the wavelength of the line of A, X is the lower-level excitation potential in eV, WA is the equivalent width of the line in rnA, and (EL/H) is the mean abundance given as the logarithmic number fraction relative to hydrogen, where H = 12. The sources of log gfvalues are as follows: [1] Becker & Butler (1990) [2] Biemont & Godefroid (1980), [3] Biemont et al. (1981), [4] Biemont et al. (1989), [5] Biemont et al. (1993), [6] Biemont, Quinet & Zeippen (1993), [7] Fuhr, Martin & Wiese (1988), [8] Grevesse et al. (1981), [9] Hannaford et al. (1982), [10] Heber (1983), [11] Kurucz & Peytremann (1975), [12] Kurucz (1988), [13] Martin, Fuhr & Wiese (1988), [14] Morton (1991), [15] McEachran & Cohen (1982), [16] Pirronello & Strazzulla (1981), [17] Sigut & Landstreet (1990), [18] Thevenin (1989), [19] Wiese & Fuhr (1975), [20] Wiese & Martin (1980), [21] Wiese, Smith & Glennon (1966), [22] Wiese, Smith & Miles (1969). lence, its origin is often considered as being mysterious. However, recent work has suggested that microturbulence is just the small-scale part of the photospheric convective flow pattern (Holweger & Sturenburg 1993). The value of ~t was obtained using Fe I lines, which dominate the optical spectrum of IK Peg A. The method of Magain (1984) was employed and the solution that gave no correlation between abundance and synthetic equivalent width was ~t = 2.6 ± 0.2 km s -1, from 104 F I lines. This result agrees well with that given in fig. 1 of Coupry & Burkhart (1992), but is lower than the values that have been associated with Am stars in the past (Smith 1973; Takeda 1984; Guthrie 1987; Ko<;er et a1. 1993). 3.3 Rotation velocity Wonnacott et a1. (1994) suggested that the value of v sin i quoted in Hoffleit (1982) was too high. Using mediumresolution spectra they gave an upper limit of v sin i ;:S50 km S-I. The availability of a high-resolution spectrum enables an accurate value of v sin i to be determined. Several of the unblended absorption lines identified in Section 3.1 were used to obtain values of v sin i. For each line the abundance required to match the measured equivalent width can be determined using UCLSYN. A synthetic spectrum is then calculated and convolved with the instrumental and rotational profiles. The value of v sin i is varied until the best fit to the observed is found. The best-fitting value is v sin i = 32.5 ± 2.5 km s -1. At the rotational velocity, the exact shape of the instrumental profile was found to be insignificant. 4 ELEMENTAL ABUNDANCES Having identified and measured the equivalent widths of the absorption lines, UCLSYN was used to calculate the elemental abundance for each line. For unblended lines this is a simple matter of determining the synthetic equivalent width that agrees with the observed value. In several cases, however, the absorption lines are significantly blended and spectrum synthesis was used to determine the abundances of the blending components. For a blended line, the abundance of the blending components are obtained from the mean values determined from single lines. The abundance of the dominant component is varied until the best-fitting least-squares solution is found. This abundance is then used to calculate the synthetic equivalent width of that line. As stated in Section 2, there is very good agreement between our Lick solar spectrum and the Kitt Peak atlas, which demonstrates that there is no significant scattered light in the observations. Hence, the equivalent widths measured in Section 3.1 should be free from any systematic errors. Another possible source of systematic error is a poor choice of log gf values. However, comparison with lines common to those of Adelman (1987) reveal no significant differences. In fact, if his equivalent widths for the F star Yf Lep are analysed using UCLSYN, we recover his abundances to within ~ 0.1 dex or better. The errors on the abundances of the individual lines were obtained from the combination of the errors due to the uncertainties in Terr, logg, ~t' loggfand W,!. A weighted mean was then used to determine the mean abundance for each element (Table 3). In some instances, abundances for two ionization stages were available. It was found that, in all cases, the difference in the means for the two ionization stages differed by less than ~ 0.1 dex. The errors on the final abundances were obtained using the procedure developed by Smith (1993). The overall abundance pattern is shown in Fig. 1. Certain groups of elements are now discussed individually. 4.1 C, Nand 0 Carbon is almost certainly slightly underabundant. This is a characteristic of Am stars (Conti 1970). Nitrogen appears to be solar, but the available optical lines are weak and noisy. The same is true of the oxygen lines, with the exception of the strong 0 I triplet around 7770 A. However, these lines are subject to NLTE effects (Faraggiana et a1. 1988). The equivalent width of the triplet is 713 rnA. If this value is divided by the correction factor given in table 3 of Faraggiana et a1. (1988), a corrected equivalent width of ~ 400 rnA is obtained. This gives an abundance of © 1996 RAS, MNRAS 278,688-696 © Royal Astronomical Society • Provided by the NASA Astrophysics Data System 1996MNRAS.278..688S The chemical composition of IK Pegasi (O/H)=8.6 or [O/H) = -0.4. In addition, inspection of table 4 of Faraggiana et al. reveals that the measured equivalent width is lower than the computed value for a solar oxygen abundance with atmospheric parameters adopted for IK Peg A. Hence, there are indications that oxygen is also slightly underabundant. Table 3. Mean elemental abundances. EL C N 0 Na Mg AI Si S K Ca Sc Ti V Cr Mn Fe Co Ni Cu Zn Sr Y Zr Ba n 18 3 3 2 3 4 16 2 1 20 7 19 8 19 9 119 3 18 2 2 4 9 7 6 (EL/H) 8.26 ± 0.18 8.03 ± 0.19 8.75 ± 0.19 6.40 ± 0.13 7.45 ± 0.23 6.30 ± 0.17 7.68 ± 0.16 7.20 ± 0.15 4.88 ± 0.14 6.25 ± 0.25 2.91 ± 0.20 5.25 ± 0.24 4.07 ± 0.16 5.80 ± 0.19 5.44 ± 0.18 7.71 ± 0.17 5.14 ± 0.15 6.54 ± 0.19 4.26 ± 0.21 4.68 ± 0.17 3.39 ± 0.17 2.32 ± 0.15 2.53 ± 0.18 3.26 ± 0.17 (EL/H)0 8.60 8.00 8.94 6.33 7.58 6.47 7.55 7.21 5.12 6.36 3.10 4.99 4.00 5.67 5.39 7.54 4.92 6.25 4.21 4.60 2.90 2.24 2.60 2.13 [EL/H] -0.34 ± 0.18 +0.03 ± 0.19 -0.19 ± 0.19 +0.07 ± 0.13 -0.13 ± 0.23 -0.17 ± 0.17 +0.13 ± 0.16 -0.01 ± 0.15 -0.24 ± 0.14 -0.11 ± 0.25 -0.19 ± 0.20 +0.26± 0.24 +0.07 ± 0.16 +0.13 ± 0.19 +0.05 ± 0.18 +0.17 ± 0.17 +0.22 ± 0.15 +0.29 ± 0.19 +0.05 ± 0.21 +0.08 ± 0.17 +0.49 ± 0.17 +0.08 ± 0.15 -0.07 ± 0.18 +1.13 ± 0.17 Notes: n is the number of lines used in the means. (ELI H) is the mean abundance given as the logarithmic number fraction relative to hydrogen, where H = 12. (ELlH)o denotes the solar elemental abundances taken from Anders & Grevesse (1989). [ELlH] is the mean logarithmic elemental abundance relative to the solar value. N No AI Lo Pr K Sc V Mn Co eu 4.2 Ca and Sc One of the defining characteristics of an Am star is weak Ca Hand K lines, which imply an underabundance of Ca. However, these lines could not be analysed because their profiles were too broad to be accurately extracted from the echelle spectra. Nevertheless, 20 Ca I lines were measured, yielding a mean abundance that was solar to within the estimated errors. Also, the Sc abundance was solar to within the estimated errors. 4.3 Fe group Another of the defining characteristics of an Am star is strong metallic lines compared with a normal star of similar spectral type. In an A-type star the metallic lines are due to the Fe-group elements and in particular Fe I. The implication of this spectroscopic definition is that the Fe-group elements are overabundant relative to the solar abundances. In typical classical Am stars, enhancements of ~ 0.5 dex are common (Conti 1970). In the case of IK Peg A, there is an overall abundance enhancement of + 0.2, which is in good agreement with the value of mean metal abundance, [M/H), obtained by Wonnacott et al. (1994). Indeed, if the method of Smalley (1993) is applied to a suitably degraded version of the high-resolution spectrum, a value of [M/H) = 0.12 ± 0.15 is obtained, which is also in very good agreement. This, in combination with the Ca result, means that IK Peg A cannot be a classical Am star according to the definition of Conti (1970). 4.4 Cu, Zo, Sr, Y, Zr and Ba Of these elements, the considerable overabundance of Ba is slightly significant. An overabundance of Ba is a characteristic of Am stars, but so is an overabundance of Cu, Zn, Sr, Y and Zr. Of these elements, only Sr is enhanced; the rest are solar. This abundance pattern is not typical of an Am star. Ell Table 4. Rare-earth elemental abundances. EL La Ce Pr Nd Sm Eu Gd Dy Yb Lu Element Figure 1. The abundance pattern for IK Peg A. The filled circles are elements with reliably determined abundances (Table 3), while the rare-earth elements are shown as open circles. With the exception of Nd, all the rare-earth abundances are upper-limits. 693 n 2 (EL/H) <2.2 <2.0 <1.1 2.14 ± 0.23 <2.0 <0.7 <1.3 <1.4 <2.4 <1.2 (EL/H)0 1.22 1.55 0.71 1.50 1.01 0.51 1.12 1.10 1.08 0.76 [EL/H] <+1.0 <+0.4 <+0.4 +0.64 ± 0.23 <+1.0 <+0.2 <+0.2 <+0.3 <+1.3 <+0.4 Notes: n is the number of lines used in the means. (ELI H) is the mean abundance given as the logarithmic number fraction relative to hydrogen, where H = 12. (ELlH)o denotes the solar elemental abundances taken from Anders & Grevesse (1989). [EL/H] is the mean logarithmic elemental abundance relative to the solar value. © 1996 RAS, MNRAS 278, 688-696 © Royal Astronomical Society • Provided by the NASA Astrophysics Data System 1996MNRAS.278..688S 694 B. Smalley et al. 4.5 Rare earths Table 4 summarizes the results of a systematic search for lines due to rare-earth elements. These elements are usually considerably enhanced (> 1 dex) in classical Am stars (Smith 1971; Hundt 1972). Only Nd has been positively identified and it is only slightly enhanced. Other rare-earth elements have upper-limits that imply that they cannot be much enhanced from the solar values. The rare-earth abundances for IK Peg A are not that of a classical Am star. 5 DISCUSSION A detailed elemental abundance analysis of IK Peg A has revealed that the Ca and Sc abundances are roughly solar, while the Fe-group elements are slightly enhanced. The ratio of [Ca/Fe] = [Ca/H] - [Fe/H] can be used as a measure of the degree of Am nature of a star. For IK Peg A, we have [Ca/Fe] = - 0.28 ± 0.30, which indicates that there is, at most, only a very mild Am character to this star. A classical Am star, such as 63 Tau, has a ratio of [Ca/Fe] = -1.0 (Burkhart & Coupry 1989). This shows conclusively that IK Peg A is not a classical Am star. The results are not, however, inconsistent with its spectroscopic classification as a marginal Am star. Given the observed overabundances, can they be satisfactorily explained? There are two possibilities, namely that the barium and strontium excesses are the product of radiative levitation acting on normal abundances of these elements, or that the overabundances are real and the elements have been transferred from the giant companion during the common envelope (CE) phase. IK Peg A is known to be a multimode pulsator: the Am and b Scuti phenomena are thought to be mutually exclusive, because radiative diffusion models imply the gravitational settling of helium upon which the mechanism for pulsation relies (Michaud et al. 1983). Calcium and scandium underabundances, if due to radiative diffusion, require the disappearance of the helium convection zone. In IK Peg A, Ca and Sc are essentially normal, implying that the helium convection zone has not disappeared, and the source of the observed pulsations is therefore available. Nevertheless, some heavy elements can become overabundant, even if there is too much turbulence (caused by rotation or pulsation for example) for the helium to settle gravitationally (Vauclair, Vauclair & Michaud 1978). Hence, the first possibility is plausible provided that the precise observed abundance pattern can be explained, i.e. why the (apparent) excess of s-process elements is seen rather than those patterns observed in classical Am stars, or HgMn stars, for instance. With appropriate modelling, it may be possible to test if the observed levels are achievable solely under radiative levitation, or if there are simply too many of the s-process elements seen to be explained in this way. Then only the second possibility remains. This requires that the s-process elements are produced by the activation of a source of neutrons as the white dwarf progenitor enters the giant branch. As the CE phase progresses and the A-star companion penetrates the giant's atmosphere more and more deeply, these newly created elements are transferred with some (unknown) efficiency to the A-star. There is, however, a problem of timing with this hypothesis. The CE phase is expected to last a very short time [hundreds of thousands of years; see Iben (1991)]. The manufacture of the s-process elements proceeds on a timescale longer than this as the activation of the neutron source takes place with the first helium flash at the tip of the RGB. The conversion of 22Ne to 25 Mg (producing the neutrons) is very efficient in high-core-mass objects such as the progenitor to IK Peg B, so the time-scale of s-process element production is that taken for say, a 5-Mo progenitor to reach this point, i.e. several hundreds of thousands of years (Iben 1991; Malaney & Lambert 1988). Of course, it may be possible that the Roche lobe surrounding the giant was not filled for some time as the latter ascended the RGB, and rapid Roche lobe overflow and the onset of the CE phase took place only within 5-10 thousand years, or less, of the red giant tip. Naturally, the exact nature of the interactions during the CE phase depends upon the initial sizes of the orbital separation and the 'final' size of the giant. If, for example, the A star was almost as far away as the giant's radius, then the AGB phase would be of maximum duration and any mass transfer would have taken place near the surface, requiring a full dredge-up of s-process elements. However, if the A-star orbit was much closer the star would have been plunged much deeper into the giant's atmosphere as it expanded and accreted s-process material more directly. Obviously, the detailed modelling needed to confirm or refute either of the above hypotheses are outside the scope of this paper. The authors merely note that either possibility is consistent with the observed facts. Finally, there is the question of the existence of objects similar to IK Pegasi. The enhancement of barium and the presence of a white dwarf companion naturally brings the barium stars to mind [discovered by Bidelman & Keenan (1957)], but the primaries of these binaries are evolved and the secondaries (the white dwarfs) are of lower masses [~0.6 Mo; McClure & Woodsworth (1990)]. Futhermore, the brevity of the period would seem to mitigate against the idea that IK Peg is a Ba II star as the latter type has periods in the range 80-2000 d or longer. Two other groups of stars also possess similarities: the CH giants and subgiants, and the F str ),4077 stars. The former group can be ruled out as they are, like the barium stars, known to be long-period binaries (North & Duquennoy 1992) and possess a C/O ratio of order unity or greater (Luck & Bond 1991); IK Peg A has C/O = 0.32 ± 0.20. The F str ).4077 stars [discovered by Bidelman (1981)] are more promising. North & Duquennoy (1991) have made a study of these objects and conclude that they can be in binaries with periods as short as a few days, that they rotate more slowly than field A-F stars, and have an excess of strontium. The F str ).4077 stars, in common with the Am and barium stars, generally tend to have excesses of Y and Zr. IK Peg A has normal abundances of these two elements and as such is somewhat atypical. Yet another point of coincidence between this group and IK Peg A is that the F str ).4077 stars overlap with an instability strip on the HR diagram. On fig. 4 of North & Duquennoy (1991), IK Peg A is located about two-tenths of a magnitude inside the red edge of this strip, nicely accounting for its pulsation. Finally, the presence of a massive white-dwarf companion may be an evolutionary requirement: a hot A-star will evolve faster © 1996 RAS, MNRAS 278, 688-696 © Royal Astronomical Society • Provided by the NASA Astrophysics Data System 1996MNRAS.278..688S The chemical composition of IK Pegasi than a cooler F-analogue and so will require a more massive companion to out-evolve it, if it is still to produce the observed system - hence, a more massive remnant core survives the stellar disruption giving rise to the massive white dwarf observed. In summary, only one of the binary CP-star groups (albeit a rather heterogeneous one) is consistent with observed abundances, pulsation, and binary composition and structure of IK Pegasi. It is posited then, that IK Peg A is a hot F str ,1.4077 star. The F str ,1.4077 stars are a heterogeneous class, with the hotter members possibly being related to Am stars (North & Duquennoy 1991; North, Berthet & Lanz 1994). Thus, it is unclear as to the origin of the anomalies in the hot F str ,1.4077, since their similarities to Am stars raises the possibility that the anomalies are due to radiative diffusion and not mass-transfer. In fact, even if there were any anomalies due to mass-transfer they may well have been masked by the effects of subsequent radiative diffusion (North, Berthet & Lanz 1994). 6 CONCLUSION A detailed abundance analysis of the pulsating A-type star IK Peg A has revealed that the Ca and Sc abundances are approximately solar and the Fe-group elements slightly enhanced. This star is not a classical Am star, but the results are not inconsistent with its spectroscopic classification as a marginal Am star. Whether this marginal Am character is due to the effects of radiative diffusion is unclear, because there is the possibility that some, or all, of the abundance anomalies may be the result of mass transfer. The obvious excess of Ba and Sr are cited as evidence for mass transfer on to IK Peg A from the white dwarf progenitor during the common envelope phase of the binary system evolution. It is suggested that IK Peg A could be a hot member of the F str ,1.4077 stars. This group is, however, rather heterogeneous, with the hotter members possibly being related to Am stars. This raises the possibility that the anomalies are due to radiative diffusion and not mass transfer. Further studies of other similar objects are urgently required to investigate whether elements mental enhancements that are the result of mass transfer can be observed in the atmospheres of A-type stars, or to confirm that these enhancements are effectively masked by the effects of radiative diffusion subsequent to an episode of mass transfer. ACKNOWLEDGMENTS The referee, Pierre North, is thanked for his helpful comments on the original manuscript. 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