Download The chemical composition of IK Pegasi

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. The authors thank Tony
Misch and all the staff at Lick Observatory for their hospitality and assistance during the two observing runs and Jon
Holtzman for his help in generating an OSF version of Lick
VISTA. Tony Lynas-Gray is thanked for making available a
machine-readable version of the Kitt Peak solar flux atlas.
This work has made use of the hardware and software provided at Keele and UCL by the PPARC Starlink project.
This work was based on observations obtained using the
Hamilton Echelle Spectrograph on the Conde Auxiliary
Telescope at Lick Observatory, University of California.
695
REFERENCES
Abt H. A, Bidelman W. P., 1969, ApJS, 6, 37
Adelman S. J., 1987, A&AS, 67, 353
Anders E., Grevesse N., 1989, Geochim. Cosmochim. Acta, 53,
197
Barstow M. A, Holberg J. B., Fleming T. A, Marsh M. C, Koester
D., Wonnacott D., 1994, MNRAS, 270, 499
Barstow M. A, Holberg J. B., Koester D., 1994, MNRAS, 270,
516
Batten A H., Fletcher J. M., MacCarthy D. G., 1989, Pub!. Dom.
Astrophys. Obs., 17
Becker S. R, Butler K., 1990, A&A, 235, 326
Bidelman W. P., 1981, AJ, 86, 553
Bidelman W. P., Keenam P. C, 1957, ApJ, 114,473
Biemont E., Godefroid M., 1980, Phys. SCT., 22, 231
Biemont E., Grevesse N., Nannaford P., Lowe R M., 1981, ApJ,
248, 867
Biemont E., Grevesse N., Faires L. M., Marsden G., Lawler J. E.,
1989, A&A, 209, 391
Biemont E., Hibbert A, Godefroid M., Vaeck N., 1993, ApJ, 412,
431
Biemont E., Quinet P., Zeippen C J., 1993, A&AS, 102,435
Burkhart C, Coupry M. F., 1989, A&A, 220, 197
Conti P. S., 1970, PASP, 82, 781
Coupry M. F., Burkhart C, 1992, A&AS, 95,41
Cowley C. R, Aikman G. C L., 1980, ApJ, 242, 684
Cowley C R, Henry R, 1979, ApJ, 233, 633
Cowley A, Cowley C, Jaschek H., Jascheck C., 1969, AJ, 74, 375
Faraggiana R, Gerbaldi M., van't Veer C, Floquet M., 1988,
A&A, 201, 259
Fuhr J. R, Martin G. A, Wiese W. L., 1988, J. Phys. Chern. Ref.
Data, 17, Supp!. 4
Grevesse N., Biemont E., Hannaford P., Lowe R M., 1981, in
Liege Astrophys. Colloq. No. 23, Upper Main Sequence CP
Stars. Institut d'Astrophysique, Universite de Liege, p. 211
Guthrie B. N. G., 1987, MNRAS, 226, 361
Hannaford P., Lowe R M., Grevasse N., Biemont E., Whaling W.,
1982,ApJ, 261, 736
Harper W. E., 1935, Pub!. Dom. Astrophys. Obs., 6, 249
Heber U., 1983, A&A, 118, 39
Hoffleit D., 1982, Bright Star Catalogue. Yale Univ. Observatory,
New Haven, Connecticut.
Holwegar H., Stiirenburg S., 1993, in Dworetsky M. M., Castelli F.,
Faraggiana R, eds, Proc. IAU Colloq. 138, ASP Conf. Ser., 44,
Peculiar Versus Normal Phenomena in A-Type and Related
Stars. Astron. Soc. Pac., San Francisco, p. 356
Hundt E., 1972, A&A, 21, 413
Iben I., 1991, ApJS, 76, 55
Ko<;er D., Adelman S. J., Bolcal <::., Hill G., 1993, in Dworetsky M.
M., Castelli F., Faraggiana R, eds, Proc. IAU Colloq. 138, ASP
Conf. Ser., 44, Peculiar Versus Normal Phenomena in A-Type
and Related Stars. Astron. Soc. Pac., San Francisco, p. 213
Kurtz D. W., 1978, ApJ, 221, 869
Kurucz R. L., 1979, ApJS, 40,1
Kurucz R L., Peytremann E., 1975, Smithson. Astrophys. Obs.
Spec. Rept., 362
Kurucz R L., Furenlid I., Brault J., Testerman L., 1984, Solar Flux
Atlas from 296 to 1300 nm, National Solar Observatory Atlas
No.1, Harvard Univ. Press
Landsman W., Simon T., Bergeron P., 1993, PASP, 105,841
Landsman W., Simon T., Bergeron P., 1995, in Koester D., Werner
K., eds, Proc. 9th European Workshop on White Dwarfs.
Springer-Verlag, New York, p. 191
Luck R. E., Bond H. E., 1991, ApJS, 77, 515
McClure RD., Woodsworth A W., 1990, ApJ, 352, 709
McEachran E. P., Cohen M., 1982, J. Quant. Spectrosc. Radiat.
© 1996 RAS, MNRAS 278, 688-696
© Royal Astronomical Society • Provided by the NASA Astrophysics Data System
1996MNRAS.278..688S
696 B. Smalley et al.
Transfer, 27, 119
Magain P., 1984, A&A, 134, 189
Malaney R. A., Lambert D. L., 1988, MNRAS, 235, 695
Martin G. A., Fuhr I. R., Wiese W. L., 1988, J. Phys. Chern. Ref.
Data, 17, Suppl. 3
Meyerdierks R., 1993, Starlink User Note No. 86
Michaud G., 1970, ApI, 160,641
Michaud G., Tarasick D., Charland Y., Pelletier c., 1983, ApJ, 269,
239
Morton D. c., 1991, ApJS, 77, 119
North P., Duquennoy A., 1991, A&A, 244, 335
North P., Duquennoy A., 1992, in Duquennoy A., Mayor M., eds,
Binaries as tracers of stellar formation. Cambridge Univ. Press,
Cambridge, p. 202
North P., Berthet S., Lanz T., 1994, A&A, 281, 775
Pirronello V., Strazzulla G., 1981, A&A, 93, 411
Pogge R. W., Goodrich R. W., Veilleux S., 1988, The VISTA
Cookbook, Lick Observatory Technical Report No. 50, Univ.
California, U.c. Regents
Roman N. G., Morgan W. W., Eggen O. I., 1948, ApI, 107, 107
Sigut T. A. A., Landstreet I. D., 1990, MNRAS, 247, 611
Smalley B., 1993, A&A, 274, 391
Smalley B., Dworetsky M. M., 1993, A&A, 271, 515
Smith K. c., 1992, PhD Thesis, Univ. London, Ch. 5
Smith K. C., 1993, A&A, 276, 393
Smith M. A., 1971, A&A, 11, 325
Smith M. A., 1973, ApJ, 182, 159
Stover R. I., 1988, in Robinson L., ed., Proc. Ninth Santa Cruz
Summer Workshop Astron. Astrophys., Instrumentation for
Ground-Based Optical Astronomy. Springer-Verlag, Berlin, p.
443
Takeda y., 1984, PASI, 36, 149
Thevenin F., 1989, A&AS, 77, 137
Titus I., Morgan W. W., 1940, ApJ, 92, 256
Vauclair G., Vauclair S., Michaud G., 1978, ApI, 223, 920
Vogt S. S., 1987, PASP, 99,1214
Wiese W. L., Fuhr J. R., 1975, I. Phys. Chern. Ref. Data, 4, No 2
Wiese W. L., Martin G. A., 1980, NSRDS-NBS 68, Part 2
Wiese W. L., Smith M. W., Glennon B. M., 1966, NSRDS-NBS 4
Vol. 1
Wiese W. L., Smith M. W., Miles B. M., 1969, NSRDS-NBS 22
Vol. 2
Wonnacott D., Kellett B. I., Stickland D. J., 1993, MNRAS, 262,
277
Wonnacott D., Kellett B. J., Smalley B., Lloyd c., 1994, MNRAS,
267,1045
© 1996 RAS, MNRAS 278, 688-696
© Royal Astronomical Society • Provided by the NASA Astrophysics Data System