The Astrophysical Journal Supplement Series, 32:651-680, 1976 December © 1976. The American Astronomical Society. All rights reserved. Printed in U.S.A. METALLICISM AND PULSATION: AN ANALYSIS OF THE DELTA DELPHINI STARS Donald W. Kurtz* University of Texas at Austin, and McDonald Observatory Received 1976 February 23; revised 1976 April 12 ABSTRACT Fine abundance analyses of eight 8 Delphini stars and one 8 Scuti star relative to four comparison standard stars are presented. Five of the 8 Delphini stars are shown to have abundances most similar to the evolved Am stars. It is argued that these abundances are different from the main-sequence Am star and Ap star abundances, and that similarities to the Ba n star abundances are coincidental. We suggest that the anomalous-abundance 8 Delphini stars are evolved metallic-line stars on the basis of their abundances, position in the (j8, Mv) plane, inferred rotation velocities, and perhaps their binary incidence. Some of the 8 Delphini stars are 8 Scuti pulsators. We argue that pulsation and metallicism are mutually exclusive among the classical Am stars but may coexist in other stars related to the classical Am stars. A preference for the diffusion-hypothesis model for the metallic-line stars is stated and supported, and the implications of the coexistence of pulsation and diffusion are discussed. Subject headings: stars: abundances — stars: 8 Scuti — stars: metallic-line — stars: pulsation i. introduction The 8 Delphini stars were defined as a class by Bidelman (1965), who designated 15 of 82 metallic-line stars as S Delphini from an objective-prism survey. The only clarification of the classification was that the 8 Delphini stars are metallic-line stars "in which the difference between the metallic-line type and the K-line type is rather small." The class was used by Cowley and Cowley (1964), who classified a star as having a spectrum "Ľke 8 Del." They (Cowley and Cowley 1965) reexamined Bidelman's Am and 8 Delphini stars using slit spectra, and changed the classification of some stars but did not elaborate on the 8 Delphini classification. Cowley (1968) states that" the metallic line spectrum [of a 8 Delphini star] resembles that of an F2 IV star but the hydrogen and ionized calcium lines are very narrow." This description is expanded in the Bright A Star Catalog (Cowley et al. 1969) in the description of 8 Del itself: "The spectrum of 8 Del shows rather narrow but equal H and K lines. Hydrogen lines are narrow; metallic hne spectrum is rich and similar to that of a late A metallic line star." It is further explained (Cowley and Crawford 1971) that AA 4173-4178 (Fe i, Y n, Fe n) and A 4150 (Zr n) are especially enhanced, whereas A 4417 (Ti n) is weak, as are the other metals (Cowley 1973). In the defining paper for the MKA system for F giants, Morgan and Abt (1972) classify 14 of the 16 8 Scuti variables listed by Danziger and Dickens (1967). Four of these stars, including 8 Del itself (F0 IVp), are noted to have peculiar spectra in which * Visiting Student, Kitt Peak National Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc., under contract with the National Science Foundation. the Ca ii H and K fines are weak for the MKA type. It is also noted that all of the 8 Scuti variables that they classify have a luminosity class brighter than class V. Malaroda (1973, 1975) uses the MKA criteria for 8 Scuti variables with peculiar spectra to classify some stars as 8 Delphini. We agree that the MKA peculiar 8 Scuti stars are, for classification purposes, 8 Delphini stars, but we believe that there is a possible confusion between the 8 Delphini and 8 Scuti classes which should be clarified. We use the 8 Scuti classification to refer to photometrically variable stars within 3 magnitudes of the main sequence, with periods between 0.5 and 5 hours, and with amplitudes generally less than a few hundredths of a magnitude (Baglin et al. 1973). They are interpreted as pulsational variables lying in the extension of the Cepheid instability strip where it crosses the main sequence between spectral types A2 and F0. We use the 8 Delphini classification to refer to stars with spectra similar to that of 8 Del itself—that is, late A and early F subgiants and giants with disparate K-line and metal-line spectral types. Delta Scuti is a photometric classification and 8 Delphini a spectroscopic classification; the two, as will be shown, cannot be used interchangeably. Following the initial calculations of Michaud (1970), Watson (1970, 1971) and Smith (1971) suggested that element diffusion could account for the anomalous abundance patterns seen in the metallic fine stars. Smith (1971, 1973a) used his extensive observational data to build a qualitative model for the Am stars in which it was suggested that element diffusion occurs in the radiative zone between the H i, He i, and the He ii ionization zones. This model explains the observed abundance anomalies, their temperature dependence, the low-temperature cutoff of the Am domain, 651 © American Astronomical Society • Provided by the NASA Astrophysics Data System 652 KURTZ Vol. 32 and the correlation between metallicism and rotation. The major drawback of the model is that the diffusion velocities are predicted to be on the order of 10~5cms_1, whereas there is little theoretical evidence indicating whether stability against turbulent or convective mixing on that velocity scale is plausible or not. Latour et al. (1975) suggest that convective overshoot from the He u ionization zone may disrupt the above radiative zone sufficiently that element diffusion may not be able to occur there. Breger (1970) showed that, in general, Am stars do not pulsate, and he hypothesized (Breger 1972) that, within the diffusion model for A stars, either (i) pulsation disrupts the extreme stability necessary for diffusion to occur to produce an Am star, or (ii) in a star in which diffusion does occur the helium sinks out of the He n ionization zone, thus damping the driving mechanism for pulsation in 8 Scuti stars. Baglin (1972) and Vauclair et al. (1974) calculate that, in a star in which diffusion occurs, helium sinks rapidly from the He n ionization zone. Several stars have been labeled pulsating Am stars, but Kurtz et al. (1976) have shown that other explanations are more plausible in each of these cases. There is at present no known exception to the exclusion between the classical Am stars and the 8 Scuti pulsators. Some of the 8 Delphini stars, however, appear to be related to both the Am stars and the S Scuti stars. As these 8 Delphini stars are subgiants and giants with Am-like spectra, one might a priori postulate that they may have evolved from Am stars. Some of the 8 Delphini stars are also large-amplitude 8 Scuti pulsators. This leads us to ask, Is there a region of the H-R diagram where pulsation and metallicism can coexist? If so, what effect does this have on the plausibility of the diffusion hypothesis as applied to the Am stars and to the S Scuti variables? What is the physical nature of the 8 Delphini stars? In this paper we will begin to answer these questions. In § II we analyze uvbyfi photometry of the S Delphini stars and compare it with uvbyfi photometry of the metallic-line stars. In § III the relationships among rotation, pulsation, and metallicism and their implications for the S Delphini stars are discussed. Sections IV, V, and VI present the abundance analyses of eight 8 Delphini stars and one S Scuti star relative to four comparison standards. Section VII mentions what is known about the binary incidence among the 8 Delphini stars, and §VIII is a discussion of the nature of the 8 Delphini stars and their significance to our understanding of the metallic-line phenomenon and pulsation in the 8 Scuti stars. In § IX we speculate on the nature of the proposed coexistence of diffusion and pulsation. In Appendix A we define the various subclassifications of the metallic-line and related stars as they are used in this paper. ii. photometry The 8 Delphini stars are a spectroscopically defined class of stars. Photometry, therefore, provides independent information about these stars which may be used to help determine whether the members of the class are astrophysically related. In discussing these stars we find it convenient to break them up into two subgroups based on their apparent visual magnitudes and the source of their spectral-type classifications. The first group we will refer to as the bright 8 Delphini stars. They are stars classified S Delphini or Fp from slit spectra by Cowley et al. (1969), Cowley and Crawford (1971), Morgan and Abt (1972), and Mala-roda (1975) and which have mv < 6.7 mag. The second group we will refer to as the faint 8 Delphini stars. They are stars originally classified as 8 Delphini by Bidelman (1965) from objective-prism plates and later reclassified by Cowley and Cowley (1965) using slit spectra. They have mv > 6.4 mag. There is some overlap in apparent visual magnitude between the two subgroups, so we reiterate that the subdivision is for convenience of discussion only, with no a priori implication about the physical nature of the members of each group. Table 1 lists the photometric indices of the uvbyfi system from Lindemann and Hauck (1973) for the bright 8 Delphini stars. Table 2 lists the indices obtained by the author for the faint 8 Delphini stars along with the classification of those stars by Bidelman (1965) and by Cowley and Cowley (1965). a) uvbyfi Photometry of the Faint Delta Delphini stars Observations were obtained on 1974 September 6 and September 8 and on 1975 February 15 with the University of Texas Volksphotometer attached to the McDonald Observatory 76 cm telescope. Each TABLE 1 Bright Delta Delphini Stars HR V b-y ml Ci |8 Reference* 421... . 5.68 0.208 0.151 0.674 2.726 4,6 1706... . 5.06 0.130 0.180 0.998 2.799 1 1974... . 6.44 0.160 0.175 0.764 2.746 1 2094... . 5.28 0.178 0.186 0.744 2.768 2 2100... : 5.88 0.116 0.218 0.952 2.789 1 2255... . 6.67 0.224 0.193 0.937 2.753 3 2557... . 5.98 0.221 0.142 1.023 2.741 3 3185... . 2.88 0.259 0.215 0.731 2.715 2,5 3228... . 6.38 0.174 0.221 0.831 2.775 3 3265... . 6.30 0.196 0.230 0.786 2.753 5 3649... . 6.34 0.204 0.171 0.630 2.733 4 4760... . 5.37 0.118 0.211 0.996 2.830 1 5017... . 4.71 0.180 0.231 0.913 2.780 5 6492... . 4.30 0.257 0.176 0.685 2.706 2 6561... . 3.54 0.152 0.203 0.890 2.790 2 7020... . 4.72 0.214 0.197 0.830 2.749 2, 5, 6 7859... . 5.03 0.254 0.263 0.656 2.724 2 7928... . 4.53 0.191 0.162 0.853 2.739 1,5 7984... . 5.08 0.108 0.209 0.897 2.840 1 8102... . 6.44 0.189 0.169 0.913 2.766 1 8322... . 2.83 0.184 0.186 0.744 2.768 1 8787... . 4.27 0.253 0.242 0.644 2.733 2 * Reference for the 8 Delphini classification: (1) Cowley et al. 1969, (2) Malaroda 1973, 1975, (3) Cowley and Crawford 1971, (4) Cowley 1973, (5) Morgan and Abt 1972, (6) Cowley and Fraquelli 1974. © American Astronomical Society • Provided by the NASA Astrophysics Data System No. 4, 1976 METALLICISM AND PULSATION 653 TABLE 2 Faint Delta Delphini Stars HD (D V (2) b-y (3) (4) Cl (5) ß (6) Spectral type* (7) Spectral type* (8) 3448 8.96 0.238 0.231 0.735 2.755 F3 V 8 Del 7119 7.53 0.195 0.221 0.787 2.757 S Del 18460 8.40 0.219 0.210 0.778 2.799 F3 V 8 Del 25515, 8.64 0.254 0.175 0.758 2.697 F3III 8 Del 30110. 7.43 0.192 0.202 0.729 2.760 8 Del 8 Del 39390 8.46 0.199 0.183 0.701 2.723 8 Del 8 Del 47606 7.29 0.130 0.208 1.000 2.813 8 Del Am 69682 6.47 0.184 0.195 0.712 2.761 FOIV 8 Del 72792 7.59 0.217 0.223 0.726 2.738 8 Del 8 Del 78388 7.56 0.230 0.168 0.714 2.709 F0 III 8 Del 81772 8.20 0.176 0.237 0.817 2.803 FOIV 8 Del 172743 7.63 0.217 0.192 0.821 2.771 FO V 8 Del 179641 7.81 0.205 0.205 0.730 2.755 FO IV 8 Del 213143 7.80 0.237 0.228 0.754 2.753 Fm 8 Del 213634 8.07 0.154 0.233 0.751 2.850 FO V 8 Del 223247 8.18 0.190 0.200 0.791 2.755 FOIV S Del * Classification in col. (7) according to Cowley and Cowley 1965; classification in col. (8) according to Bidelman 1965. observation consisted of four consecutive 10 s integrations in each filter in the sequence pjlwybvuuvbypwl3n, giving a total of 80 s integration time in each filter. Thirty uvby standards (Crawford and Barnes 1970) and 30 H/S standards (Crawford et al. 1966) were observed in the same manner. Extinction coefficients were determined on 1974 September 6 and applied to both the 1974 September 6 and September 8 observations. Mean McDonald Observatory extinction coefficients were applied to the 1975 February 15 data. Transformation of the program stars to the standard system was done using linear relations for y, b — y, and jS, while a color term was included in the m1 and cy transformations as given by Crawford and Barnes (1970). The mean errors (in mag) per star determined from the standard stars for all three nights are + * « \ * • N< xx • • ■ % + - x ■ m' Jr * ; x. \ * — - ■ ■ *\ L x X X + x\ X • Am + x\. X ■ A-FI x x Bright S Del _ + Faint 8 Del . x .10 b .20 .25 b-y Fig. 2.—Comparison of the )S and b — y temperature indicators for the 8 Del stars, and for a group of A-F III stars and a group of Am stars taken from Cowley et al. (1969). The solid line represents the calibrated (Crawford 1975) relation b — y = 2.943 — /3 without the Stmi or Sci terms. HR 2100 is a known spectroscopic binary (Nadeau 1952), but nothing is known about the binary nature of HR 1974. Since HjS does not suffer from line blocking as severely as does b — y, nor from reddening, we have chosen to discuss the 8 Delphini stars further in terms of the Hj8 index, which should be a more meaningful temperature indicator for these stars than b — y. c) Metallicity Index The description of S Del itself (Cowley et al. 1969)— " metallic line spectrum is rich and similar to that of a late A metallic line star"—makes one suspect that the metallicity index, m1} might be enhanced in the 8 Delphini stars as it is in many of the Am stars. Figure 3 is a plot of m1 versus j8 for the 8 Delphini stars and for a random sample of Am stars selected from the catalog of Lindemann and Hauck (1973). The mean relations for the field-star and Hyades main sequence have been drawn in with error bars encompassing 75% of the sample used to define the relation. About two-thirds of the 8 Delphini stars he within the scatter of the field-star main-sequence relation, but the other one-third do appear to have the high ntj, index indicative of increased metal line blocking, with several of them lying in the Am domain in this plot. Two of the stars, HR 7859 and HR 8787, classified 8 Delphini by Malaroda (1975), have very large mx indices and are near the cool border of the Am domain. They have been previously classified as g?F5 and F6 IV, respectively (Hoffleit 1964). As they are southern stars, we have not yet observed them, but we give them special notice here for their interesting position in the (j8, plane and their spectral type. Some caution must be used in inferring metallicity from the m1 index. First, some of the metallic-line stars do not have abnormally high mx indices (Milton and Conti 1968). A few of these are plotted in Figure 3. Second, the main-sequence relations may not apply to giant stars (Hauck 1971; Baglin et al. 1973). For some giants, m1 decreases relative to the main-sequence value at the same j8, so that a giant with an increased metal abundance may have an mx index very near the main-sequence value at that /3. The high mx indices of some of the S Delphini stars very probably imply a high metallicity in those stars, but the normal mi indices of the rest do not necessarily imply a normal metallicity. d) Position in the (fi, Mv) Plane Using Crawford's (1970) calibration of Mv for A stars, we have plotted in the (j8, Mv) plane the faint Fig. 3.—The metallicity index, mu in the S Del stars compared with a sample of Am stars selected from the catalog of Lindemann and Hauck (1973). The solid lines represent the mean mu p relation for field stars (Crawford 1970) and for the Hyades (Breger 1968). The error bars are drawn to include 75% of the sample from which the mean relations were derived. © American Astronomical Society • Provided by the NASA Astrophysics Data System No. 4, 1976 METALLICISM AND PULSATION 655 i.o Mv 2.0 ... M 1 1 1 1 1 1__ 1 1 1 X 1 1 1 1 "'■ ! 1 l / / / / / / / 1 1 _ 1 1 —/ • x / • _ 1 A2^----• ZAMS "^""^^^^^ x!/ FO 1 i i i ■ Iii 2.82 2.78 2.74 2.70 ß Fig. 4.—The faint 8 Del stars. Crosses represent stars classified 8 Del by both Bidelman (1965) and Cowley and Cowley (1965). The dashed lines delineate the observed instability strip (Baglin et al. 1973). The luminosity class IV mean relation is taken from Allen (1963). M. 2.86 2.82 2.78 2.74 2.70 yS Fig. 5.—The bright 8 Del stars, indicating the stars which have been tested for light variability. The dashed and solid lines are the same as in Fig. 4. 8 Delphini stars in Figure 4 and the bright 8 Delphini stars in Figure 5. The ZAMS is as given by Crawford (1970), while the dashed lines represent the observed boundaries of the instability strip near the main sequence (Baglin et al. 1973). The error bars represent the internal uncertainty in the Mv calibration of ± 0.3 mag. The lines of luminosity class are taken from Allen (1963). The faint 8 Delphini stars in Figure 4 seem, with a few exceptions, to be a homogeneous group of late A and early F main-sequence and subgiant stars. The crosses represent the faint S Delphini stars that were so classified by both Bidelman (1965) and Cowley and Cowley (1965). The bright 8 Delphini stars in Figure 5 show much more scatter. Most appear to be A subgiant and giant stars, while several are F main-sequence stars and one, HR 2557, is a bright, luminosity-class III star. e) Variability Some of the 8 Delphini stars show periodic light variability characteristic of pulsation, and are therefore members of the 8 Scuti class of variable stars. Others, however, are constant to less than a few thousandths of a magnitude, while still others lie outside the observed instability strip as shown in Figure 5. Table 3 is a listing of the data on all of the 8 Delphini stars which have been tested for light variability. The 8 Delphini stars which are 8 Scuti variables all have relatively large amplitudes. Three of them, p Pup (HR 3185), 8 Del itself (HR 7928), and 8 Set itself (HR 7020), are members of the original five stars from which Eggen (1956) announced the existence of the 8 Scuti class. TABLE 3 Variability among Delta Delphini Stars HR P (day) Amplitude (mag) Constancy mag hr Source* v sin j Source* 421......... Var.? 2 1706......... 0.087 o!o80 1 33 Y 1974......... Const. o!oo3 2.6 3 80 4 2100......... 0.060 1 70 1 3185......... 0.141 o'.iöo 1 14 1 3228......... Const. o!6Ö4 3.6 3 80 4 3265......... 0.097 6.Ö40 1 25 1 4760......... Const. 0'.002 4.6 3 93 4 5017......... 0.135 6.35 1 17 1 7020......... 0.194 0.290 1 32 1 7928......... 0.153 0.050 1 41,25 1,5 7984......... Const. o!ÖÖ2 L3 3 90 * (1) Baglin et al. 1973, (2) Kukarkin, Efremov, and Kholopov 1958, (3) Breger (private communication), (4) Danziger and Faber 1972, (5) personal estimate. American Astronomical Society • Provided by the NASA Astrophysics Data System 656 KURTZ Vol. 32 /) Summary and Discussion of the Photometry The preceding sections have shown that the stars classified spectroscopically as S Delphini stars are not an astrophysically homogeneous group on the basis of uvbyfi photometry. We have suggested that the computed color excesses in most of the 8 Delphini stars can be attributed to reddening, line blocking due to increased metaUicity, uncalibrated luminosity effects, or a combination of these. We have chosen j8 as a temperature parameter for the 8 Delphini stars because it is relatively insensitive to reddening and line blocking and because, for luminosity class IV and V A stars, temperature is a single-value function of P with very little dependence on luminosity (Breger 1974Z>). The mx index indicates that some of the S Delphini stars are metal-rich compared with the field star or, in some cases, even the Hyades main sequence, but that most 8 Delphini stars have mx indices within the range of normal-abundance stars. We cannot conclude, however, that these 8 Delphini stars with normal mt indices have normal abundances, since about half of the Am stars have mx < 0.230 (Milton and Conti 1968; Conti 1970), which is also within the range of normal stars. Assuming normal masses, the calculated absolute magnitudes of the 8 Delphini stars indicate that they range in luminosity from class V to class III. The faint S Delphini stars are a much more compact group in the ()3, Mv) plane than are the bright 8 Delphini stars, but many of the faint 8 Delphini stars have been reclassified as normal by Cowley and Cowley (1965). Half of the bright 8 Delphini stars have been tested for fight variability, and six of these are known to be pulsational variables of the 8 Scuti class. Several, however, are photometrically constant to a few thousandths of a magnitude, while others lie outside the present observed cool boundary of the instability strip. Because of the inhomogeneity indicated by uvbyj3 photometry of the stars comprising the 8 Delphini class, it is unsafe to assume physical parameters for a particular star of this class based only on its spectral similarity to 8 Del itself. It is certainly incorrect to assume that all 8 Delphini stars are 8 Scuti pulsators. iii. rotation, pulsation, and metallicism Before discussing rotation in the 8 Delphini stars, it is necessary to discuss the relations among rotation, pulsation, and metallicism in general for stars in the same region of the H-R diagram as the 8 Delphini stars. Figure 6 is a plot of v sin i versus amplitude of the light variability for all of the 8 Scuti stars for which both of these quantities were listed by Baglin et al. (1973). The diagram shows a clear correlation between v sin i and amplitude: the largest amplitude S Scuti stars all have low v sin i, while the fastest rotators all have relatively small pulsational amplitudes. It seems that slow rotation favors pulsation among the 8 Scuti stars. Danziger and Faber (1972) have shown that, among the late A and early F subgiant and giant stars, the slow rotators are preferentially 8 Scuti pulsators, whereas Abt (1975) has shown that virtually all of the slowly rotating A5-A9 main-sequence stars are metallic-lined. Some of the Am stars must therefore evolve into 8 Scuti pulsators, and the 8 Delphini stars are good candidates to have done just that. The mean rotational velocity of the bright 8 Delphini stars is 53 km s_1, although there is probably a selec- 160 140 120 100 80 60 40 20 ' HR3265 HRII4 • .HR7928 HR7020 I CC And* .02 .04 .06 .08 .10 .12 .14 Amplitude (mag) .20 .22 .24 Fig. 6.—The rotational velocity versus amplitude of pulsation plot for the S Scuti stars listed by Baglin et al. (1973). The correlation between these two parameters indicates that low rotational velocity is conducive to large amplitude pulsation. The low v sin ( stars are labeled. Half of them are listed in Table 1 as <5 Del stars. © American Astronomical Society • Provided by the NASA Astrophysics Data System No. 4, 1976 METALLICISM AND PULSATION 657 tion effect in favor of slow rotators, as the line strength anomalies which characterize these stars are more difficult to recognize at higher rotational velocities. In Figure 6 the 8 Scuti stars with v sin i < 40 km s_1 have been identified, and a check of Table 3 shows that one-half of them are classified as S Delphini stars. The correlation between v sin i and amplitude suggests that most of these six S Delphini stars may be intrinsically slow rotators. Assuming normal masses and evolutionary tracks, they were probably Am stars when they were on the main sequence. That the spectra of the 8 Delphini stars resemble the spectra of the classical Am stars at classification dispersion strengthens the suspicion that these pulsating subgiant and giant stars may have evolved from nonpulsating metallic-line stars. We have performed differential abundance analyses on many of the stars labeled in Figure 6 in order to test the hypothesis that such an evolutionary relation exists. iv. abundances of the sharp-lined, bright delta delphini stars Table 4 is a list of the stars on which we have performed abundance analyses along with the adopted atmospheric parameters, rotational velocities, and classification information. Preliminary differential abundances for HR 6561 were kindly provided by Myron Smith prior to publication. a) The Derived Atmospheric Parameters Effective temperatures and gravities were initially derived using Breger's (1974a, b) calibration of the (S, b — y, and Ci indices. In the case of stars for which /3 and b — y did not agree, a weighted mean was used, with |3 generally preferred. These temperatures and gravities were checked by fitting Edmonds, Schlutter, Wells (1967) theoretical Hy profiles for many of the program stars. Good agreement was obtained. In the worst case, p Pup, where p and b — y disagree, the effective temperature difference for the two parameters is about 250 K. In the cases of S Del itself and 20 CVn, which have been analyzed by other investigators, we find that the effective temperatures derived from continuum scans, fitting of Balmer fine profiles, uvbyfi photometry, and excitation equilibria, have a range of 300 K. From this we estimate the internal accuracy of our temperatures to be ±150 K. Under the assumption that the photometric temperature was correct, the surface gravity was adjusted to balance the Fe ionization equilibrium to within 0.1 dex. The surface gravities derived in this manner are systematically +0.08 ± 0.20 dex larger than the photometric surface gravities. We estimate our internal error in the surface gravity to be ±0.2 dex. Since many of these stars are S Scuti pulsators, they have variable effective temperatures and surface gravities. Baglin et al. (1973) indicate that the larger amplitude S Scuti variables have effective temperature variations of AreffXJl00K and surface gravity variations of A log g x 0.1 dex. These variations are sufficiently small that in these abundance analyses all stars will be treated as if their atmospheres were in a steady state. The microturbulent parameter, $t, was derived by requiring that no correlation exist between the derived abundances and the equivalent widths for the Fe i lines. Values in the range 4.5 < £t < 7.0 km s"1 were obtained with an estimated internal accuracy of ±0.5 kms-1. While these numbers are typical of the values derived for microturbulence by model-atmosphere abundance analyses of late A stars, systematic effects contribute considerably to the derived value. Smith (1973ft) estimates that the neglect of line blanketing, the use of the Corliss-Warner oscillator strengths, low values of the damping constants, and the large equivalent width scale associated with 8 A mm-1 plate material as compared with the equivalent widths derived from higher dispersion plate material, all contribute about 3-3.5 kms-1 to the derived microturbulent parameter. He finds further corroboration for this result (Smith 1976) by a TABLE 4 Atmospheric Data for Program Stars and Comparison Standards Teff & v sin / Star Name HR (K) (cgs) (kms-1) (kms"1) Spectral Type Source* 44 Tau....... 1287 7150 3.4 5.0 20 F2IV-V 2 14 Aur....... 1706 7900 3.8 5.0 33 8 Del 1 6 Mon....... 2255 7500 3.6 5.0 <10 s Del 4 2557 7400 3.4 5.5 30 8 Del 4 P Pup........ 3185 7100 3.25 6.0 14 F5Hp 2 3265 7450 3.75 7.0 25 F3nip 2 20 CVn 5017 7500 3.7 5.0 17 F3 HI 2 f Ser........ 6561 8000 3.9 6.8 36 8 Del 3 SDel........ 7928 7320 3.25 4.5 20 8 Del 1 28 And 114 7500 3.5 5.5 22 A7 HI, F2 V 6,7 y Vir........ 4825 7100 4.3 5.0 27 F0 V 5 8120 7600 2.5 5.0 20 Foni 1 8272 7840 3.35 6.0 20 A7 III 1 * Sources of the spectral types are: (1) Cowley et al. 1969; (2) Morgan and Abt 1972, (3) Malaroda 1975, (4) Cowley and Crawford 1971, (5) Hoffleit 1964, (6) Cowley and Fraquelli 1974, and (7) Conti 1970. American Astronomical Society • Provided by the NASA Astrophysics Data System 658 KURTZ Fourier analysis of a line profile in the Am star 32 Aqr which yields gt = 3kms_1, as opposed to it = 9kms_1 from his (Smith 1971) model-atmosphere analysis of the same star. As our analysis is quite similar to Smith's, we consider the same systematic effects to be present in our derived micro-turbulent parameters. b) Data Acquisition and Reduction For each comparison and program star, one or two Ila-O plates of reciprocal dispersion of 8 to 10 A mm-1, projected slit width of 20 fim, and widening of 0.4 to 0.8 mm were used to obtain equivalent widths. About half of the plates used were obtained by the author with the McDonald Observatory 2.7 m and 2.1 m telescopes. The rest were generously taken at the KPNO, Lick, Mount Wilson, and McDonald Observatories by Myron Smith, Michel Breger, Leonard Kuhi, Deane Petersen, and Frank Fekel. All plates were traced on the KPNO PDS micro-densitometer and converted to intensities using the KPNO spectrophotometric reduction programs, SPECT1 and SPECT2, on the University of Texas CDC 6600 computer. Equivalent widths were measured treating all lines as triangles. A complete list of all the plate material used, its source, the oscillator strengths used, and the measured equivalent widths can be found in Appendix B. In addition, the derived equivalent width scales are discussed and compared internally for material from different telescopes and externally with other published equivalent widths. Abundances were computed using convective, metal-line-unblanketed ATLAS5 (Kurucz 1970) model atmospheres with solar abundances in conjunction with the program WIDTH5 (Kurucz, private communication). Damping constants for all lines were presumed to be 10 times the value of the classical radiation damping constant. No lines on the damping portion of the curve of growth were used. We assume that errors due to incorrect oscillator strengths or damping constants are systematic and hence approximately cancel out in the differential analysis of the program stars relative to the comparison standards. We have intentionally used unblanketed models and the "old" oscillator strength scale to keep our derived abundances as close as possible to the system of Smith (1971, 1973a). Due to the similarity between our analysis and Smith's, the variation of abundance with temperature, gravity, and microturbulence computed and tabulated by Smith (1971) may be applied to the abundances presented in this paper. c) The Derived Abundances Table 5 is a listing of the log of the derived abundances on a scale of log H = 12.00, the number of lines of each ion measured, and the rms scatter. We estimate the internal error associated with these abundances to be ±0.1 dex for ions with many lines, such as Fe I, Fe II, and Ti n. This corresponds to the change in abundance for Fe if the effective temperature is changed by the estimated error of 150 K, with the necessary adjustment of the surface gravity to rebalance the Fe ionization equilibrium. The error in the abundance of ions for which only one or two lines were measured is estimated to be as large as ±0.5 dex, especially for abundances determined exclusively from very weak lines, such as Nd, or exclusively from strong lines lying on the flat portion of the curve of growth, such as Sr n and Ba n. The abundances of C i, S i, and Zn i have been determined from a few lines in the wavelength region 4700-4800 A where the IIa-0 plate is dropping in sensitivity and hence are less reliable than many of the other abundances. The rms scatter listed in Table 5 is not considered to be a good representation of the errors in the associated abundances as it includes the systematic effects of the errors in the oscillator strengths. In Table 6 we compare the derived abundances for the comparison standards HR 114 and HR 4825 with the abundances derived for the same stars by Smith (1971). We find an average scatter between the two studies of ±0.16 dex. While some of this scatter is intrinsic, part of it is systematic and due to differences in effective temperature, surface gravity, and fine lists used. The equivalent width scales for HR 114 from this study and Smith's study show no systematic shift, while for HR 4825 our equivalent widths are on the average 11% larger than Smith's. In the following discussion we choose to analyze the derived abundances normalized to the Fe abundance. This will allow an easy comparison with the metallic-line star abundances, which are best represented in this form. In addition, small shifts in the effective temperature or equivalent width scale in a given star to first approximation give rise to a shift in the entire abundance scale. Because of this the normalization to Fe minimizes the effects of errors in equivalent width scale or effective temperature. As the analyzed S Delphini stars did not all prove to have similar abundances, we will break the discussion of them into two parts. First, we will discuss the standard star abundances and will discuss as a group HR 1706, HR 2255, HR 3265, HR 6561, and HR 7928, which have similar abundance anomalies and which will hereafter be referred to as a group as the anomalous-abundance 8 Delphini stars. Then we will discuss the other analyzed stars on an individual basis. v. the abundance anomalies in HR 1706, HR 2255, HR 3265, HR 6561, and HR 7928 In Figure 7 the abundances normalized to Fe for HR 1706, HR 2255, HR 3265, HR 6561, and HR 7928 and the four comparison standards HR 114, HR 4825, HR 8120, and HR 8272 are plotted. The rms scatter for all ions for the relative abundances among the standard stars is ±0.12 dex. We find no significant systematic effects in the abundances of the standard stars as a function of luminosity. The internal scatter in the standard star abundances is less than the external differences found in the previous section in the comparison of the abundances of HR 114 and HR 4825 with those derived by Smith (1971). © American Astronomical Society • Provided by the NASA Astrophysics Data System 1976ApJS...32..651K ® > IT 2_ f?" s > v. O C 3 r? SL o 13 -8 O a a sr IT > o sr r?" Ö ta TABLE 5 Log of the Derived Abundances for Comparison and Program Stars (log H = 12.00) HR114 HR4825 HR8120 HR8272 HR1287 HR1706 HR2255 HR2557 HR3185 HR3265 HR5017 HR7928 c I 8 18( 2) .20 7. 89( 2) .09 7 91( 3). 09 8, 13( 3) ,24 - 8. 16( 3) ,24 C I 8 38( 2) .01 8 ,05( 2) .05 7 96( 2) .12 8. 31( 2) .12 8. ,49( 2) .15 8. 14( 3) .12 Al I 5, .08( 2) .04 4, 43( 2) ,00 4, ,66< 2). 16 5, 19( 2), ,07 4. 80( 2), ,20 5, 05 ( 2) ,29 AI I 5 45( 2) .02 4 ,78( 2) .36 5 06( 2) ,25 4, 83( 2) .37 5. ,49( 2) .26 4. 97( 2) .15 Si II - 7, .05( 2) .23 7 ,28( 2). 37 7. 06 ( 1) - 8. ,16( 1) - 7. 59 ( 2) .41 Si II 8 ,02( 2) .28 7 •14( 1) - 8 ,65( 2) .39 7 ,80( 1) - 8 53( 2) .17 7, 58( 2) .18 S I 6 56( 2) .22 6, 70( 2) .05 6, 45( 2). 17 6, 99( 3) ,13 - 6, 89 ( 3) .09 S I 7 ,00( 3) .08 7 ,06( 3) .13 5 45( 3) .32 6 50( 3) .07 7 03( 3) ,06 6. ,83( 3) .18 Ca I 5 95( 8) .19 5, 59( 7) .24 5 50( 8). 28 6, 01( 8) .28 6, 07( 9) .22 5. 99 ( 8) .26 Ca I 6 ,20( 8) .28 5 • 91( 6) .28 6 21( 8) .25 5, 94( 7) .25 6 47( 7) .29 5. 69( 8) .29 Sc II 2. 84( 7) .28 2. 31( 6) .12 2, 68( 7). 25 2. 99( . 8), .17 2. 72 ( 8) .24 3. U( 7) .27 Sc II 3. ,32( 9) .31 3 ■ 43( 8) .48 2. .92( 9) .33 2, ,82( 6). .27 3. 64( 9). .30 2. 31( 9) .21 Ti II 4. 19( 30) .29 3. ,80( 35) .25 4. ,05( 34).22 4, 41( 34), .22 4. 39( 35) .26 4. 43( 38) .29 Ti II 4, ,61( 37) .31 4 • 45( 30) .30 4, .51( 33) .30 4, ,20( 33) .20 4. ,92( 38) .28 3. 95( 36) .28 V II 3 33( 6) .25 3, 34( 6) .34 3 36( 5). 18 3, 65( 4) ,27 3, 50( 6) ,14 3, 96( 6) .23 V II 4 OK 7) .24 3 ,80( 5) .34 3 78( 5) .21 3 74( 4) .20 4 26( 6) .30 3. 19( 7) .18 Cr I 4, 97( 13) .16 4. 62( U) .25 4. 71( 8). 13 5. ,15< 6) .13 5, 20( 7) .20 5. 12( 12) .22 Cr I 5, ,30( 12) .22 5 .19( 5) .22 5 36( 11) .09 5, 09( 8) .30 5 62( 11) .15 4. 67( U) .16 Cr II 5, u< 14) .30 4. 89 ( 12) .23 4. 99( 10). 12 5, 33( 12) .13 5, 38( 12) .15 5. 25( 13) .18 Cr II 5, 50( 13) .18 5 .49( 9) .25 5 ,59( 14) .28 5, ,35( 13) .30 5 ,83( 12) .17 5, 07( 13) .20 Mi I 4. 40 ( 10) .29 4. 38( U) .38 4 ,40( 9).44 4 ,99( 10) .25 4, 82( 10) ,25 4 79( 11) .37 Ml I 5 09( 11) .29 4 ,77( 11) .30 5, .16( 10) .27 4 ,89( U) .31 5 ,34( 11) .41 4, ,61( U) .38 Fe I 6. 21(117) .25 6. 01(131) .34 6 ,05( 85).32 6 42( 96) .25 6, .38.(127) .21 6, ,45(140) .24 Fe I 6, 69(151) .25 6 .46(119) ,37 6 ,72(126) .26 6 .49(117) .26 6 ,98(154) .26 6. .13(146) .24 Fe II 6. 27( 21) .20 5. 94( 24) .23 6 13( 23).18 6 ,39( 26) .19 6 ,38( 22) .21 6. ,37( 27) .20 Fe II 6, 69( 28) .24 6 ,50( 21) ,38 6 72( 29) .27 6 ,42( 25) .26 6 97( 27) .26 6, 24( 26) .16 Co I 4. 50( 3) .50 4. 38 ( 3) .73 4 ,48( 2). 59 4 ,85( 3) .68 4, 89( 3) .74 4, 94( 3) .64 Co I 5, 19( 3) .58 4 ,84( 3) .70 5 14( 3) .66 5 03( 3) .45 5 ,49( 3) .72 4, 52( 3) .57 Ni I 5. 27( 5) .24 4. 78( 7) .25 4 .84( 7). 26 5, .08( 6) .13 5, ,06( 3) .26 5 • 17( 6) .30 Ni I 5 ,64( 8) .38 5 ,10( 5) .29 5 • 72( 7) .34 5 ,38( 8) .25 5 ,81( 7) .27 4, ,98( 6) .18 Ni II 5. 05( 2) .05 5 26( 4) .35 4 ,76( 3). 10 5. ,32( 3) .18 5. ,07( 4) .24 5, ,18( 4) .23 Ni II 5, ,65( 4) .28 5 ,36( 3) .13 5 42( 3) .17 5 ,69( 4) .36 5 ,87( 4) .21 5, .10( 4) .23 ai I 2. 76( 2) .22 2 75( 2) .13 2 ,58( 2). 22 3 28( 2) .05 - 3 04( 2) .38 Zh I 3, 63( 2) .02 2 51( 1) - 3 61( 2) .08 3 35( 2) .06 3 ,58( 2) .06 3, .23( 2) .05 Sr II 2 87< 2) .12 2, 81( 2) .12 - 3 12( 2) ,15 3, ,29( 2) .02 3 64( 2) .16 Sr II 3 85 ( 2) .08 3 68( 2) .12 3 47( 2) .20 3 ,98( 2) .23 3 70( 2) .13 3. 51( 2) .12 Y II 2, 23( 6) .18 2. 32( 5) .48 2 ,29( 4). 14 2 ,40< 6) .19 2 ,33( 7) .37 2 82( 5) .33 Y II 3, ,30( 7) .30 2 80( 4) .55 2 89( 7) .55 3, ,69( 3) .09 3 .20( 7) .38 2. 73( 7) .35 Zr II 2. 61( 5) .32 2. • 42( 5) .25 2 47( 6). 16 2 89( 5) .15 2 79( 6) .30 3 ,13( 6) .25 Zr II 3, 63( 6) .29 2 98( 4) .14 3 21( 5) .31 3 ,33( 5) .19 3 ,45( 5) .26 2. ,93( 5) .20 Ba II 1. 69( 2) .08 - 1 43( 2). 29 1 71( 1) - 1 46< 1) - 2 01( 2) .29 Ba II 2, 75( 2) .30 1 ,14( 1) - 1 80( 2) .10 1 ,96( 1) - 2 35( 1) - 2, ,27( 2) .18 La II 1. 55( 5) .37 1. 48( 5) .23 1 ,70( 6). 47 1 94( 4) .58 1 55( 5) .15 2 15( 7) .33 La II 2. 62( 8) .25 2 33( 5) ,46 2 ,04( 7) .21 2 ,36( 8) .21 2 73( 8) .22 1, ,82( 7) .04 Ce II 1. 90( 6) .15 1. 85( 6) .37 1 ,70( 6). 43 2 24( 4) .75 1 ,83( 6) .28 2 38( 7) .25 Ce II 2, 69( 7) .27 2 ,31( 7) .35 2 ,36( 7) .32 2 ,76( 7) ,25 2 ,87( 7) .34 1, ,99( 7) .23 Nd II 2. 20( 2) .50 1. 82( 2) .82 - 2 29( 2) .34 1. 84( 2) .07 2 ,56( 2) .32 Nd II 2. 58( 2) .17 2 97( 1) - 2 ,46( 2) .22 2 ,84( 2) .42 2 ,78( 2) .15 2, ,25( 2) .27 an II 1 52( 1) - 1. 51( 2) .11 1 43( 1) - 1. 98( 1) - 1. 69( 2) .21 2 ,25( 2) .01 an II 2, 36( 2) .05 2 ,03( 2) ,07 2 ,09( 2) .10 2 ,32( 2) ,18 2 43( 2) .14 1, 74( 2) .11 Eu II 0, 92( 2) .16 0 73 ( 2) ,12 - 0, 80( 1) - 0, 97( 2) .14 1. 36( 2) .20 Eu II 2, 24( 2) .23 1 ,44( 2) ,67 2 ,05( 2) .26 2 ,09( 2) .40 2 42( 2) .28 1, ,58( 2) .17 Gd II - 1. 12( 1) - 1. ,37( 1) - 1 ,81( 1) - 1 .56( 1) - 2 ,09( 1) - Gd II 2. 37( 1) - 2, ■ 14( 1) - 2 ,09( 1) - 2 • 14( 1) - 2 • 41( 1) - 1, ,70( 1) - The format of each entry in this table is the average of the log abundance derived for each ion, the number of lines measured for that ion in parentheses, and the rms scatter of the abundances. 660 KURTZ Vol. 32 TABLE 6 Comparison of the Derived Abundances for the Standard Stars HR 114 and HR 4825 from This Study and the Study of Smith (1971) HR 114 HR4825 THIS STUDY SMITH(1971a) THIS STUDY SMITH(1971a) AVERAGE DIFFERENCE* c I 8. 18 8. .37 7. .89 8. 25 -0. 28 Al I 5. 08 4. 86 4. 43 4. 85 -0. 10 Si II - 7. .00 7. .05 7. 13 (-0. 08) S I 6. 56 6. .60 6. .70 6. 61 +0. .03 Ca I 5. 95 6. 09 5. .59 5. 81 -0. 18 Sc II 2. 84 2. .73 2. .31 2. 29 +0. .07 Ti II 4. 19 4. ,24 3. .80 3. 92 -0, .09 V II 3. 33 3. .12 3. .34 3. 17 +0. .19 Cr I 4. 97 5. .21 4. .62 4. 79 -0. .21 Cr II 5. 11 5, .43 4. .89 5. 32 -0, ,38 Mn I 4. 40 4. .89 4. .38 4. 54 -0. .33 Fe I 6. 21 6. .36 6. .01 5. 99 -0, ,07 Fe II 6. 27 6. .23 5. .94 6. 06 -0. .04 Co I 4. 50 4. .90 4. .38 4. .52 -0. .27 Ni I 5. 27 5. .22 4. .78 5. 01 -0. .09 Ni II 5. 05 4. .96 5. .26 5. .14 +0. .11 Zn I 2. .76 2. .84 2. .75 2. 81 -0 .07 Sr II 2. 87 2. .72 2. .81 2. .57 +0 .20 Y II 2. 23 2. .34 2. .32 2. .03 +0 .09 Zr II 2. .61 3. .00 2. .42 2. .55 -0. .26 Ba II 1. 69 1 .92 - 1, .52 (-0. .23) La II 1. 55 1. .66 1. .48 1. .39 -0. .01 Ce II 1. .90 1. .83 1. .84 1. .77 +0 .07 Nd II 2. 20 1 .69 1 .82 1. .54 +0 .40 Sm II 1. 52 1 .40 1. .51 1. ,47 +0 .08 Eu II 0, .92 1 .27 0 .73 0. .84 -0 .23 Gd II - 1 .39 1. .12 1 .09 (+0 .03) WK> 7500 7700 7100 7100 Log g(cgs) 3 .5 3 .5 4 .3 4 .0 £t(km/sec) 5 .5 5 .0 5 .0 5 .0 * The average difference is defined as the value from this study minus the value from the study of Smith. The 8 Delphini stars plotted in Figure 7 show significant abundance anomalies (by anomaly we mean an abnormal [N/Fe] ratio). Their Fe abundances range from normal in HR 7928 to an enhancement of +0.5 dex in HR 2255 and HR 6561. The average Fe enhancement is +0.3 dex. The apparent Si ii enhancement is not considered to be significant, as the Si n abundance was determined for each star from only one or two (AA4128, 4130) partially blended, high-excitation (v = 9.8 eV) lines, and as a consequence is unreliable. All of the abundances beyond the iron-peak elements are enhanced, although the derived Eu abundance is probably too large, as we have not accounted for the effect of hyperfine splitting of the Eu fines, which acts as a pseudomicroturbulence (Hartoog, Cowley, and Adelman 1974). This effect is also unaccounted for in the metallic-line star abundances which we will compare with the program 8 Delphini stars. In Figure 8 the anomalous abundances of the five program 8 Delphini stars are compared with the abundances of the main-sequence Am stars (Smith 1971) and with the abundances of five evolved Am stars, HR 1103, HR 1248, HR 5752, HR 6559, and HR 7653, lying from 1 to 2 magnitudes above the main sequence (Smith, private communication). The run of abundances for these anomalous-abundance 8 Delphini stars is remarkably similar to the abundances of the evolved Am stars. The elements typically deficient with respect to Fe in Am stars, C, Ca, and Sc, are normal or only slightly deficient with respect to Fe in this group of 8 Delphini stars, but we note that the evolved Am stars have on the average a less pronounced deficiency of these elements than do the main-sequence Am stars. This moderating of the deficient elements has been noted before in the case of Ca (Abt 1965; Smith 1971, 1973c). The rare earths are enhanced in these 8 Delphini stars, although not quite as much as in the metallic-line stars. The [Sr/Fe] and [Y/Fe] abundances are quite similar to the Am stars, while the [Zr/Fe] abundance for these 8 Delphini stars is similar to the evolved Am stars, © American Astronomical Society • Provided by the NASA Astrophysics Data System No. 4, 1976 METALLICISM AND PULSATION 661 +1.2 + 1.0 + .8 + .6 -+ .4 + .2 0 Ä-.4 + .4 + .2 0 - 2 -.4 —I-1-1-1-r S Del Stars I I—I—I—I—r + -I—I—I—r + X + ■ 5 *. ■ + T ' xf ■ A. • HRI706 + HR3265 X HR2255 ■ HR656I A HR7928 Standard Stars t + * +. •. r*; • * j * • ***. * * + •. • HR 114 and HR4825 + HR8I20 and HR8272 3 5 § 55 S « H « 2 |5 1=1 1=1 3 3 e Fig. 7.—The derived abundances normalized to Fe in the anomalous-abundance 8 Delphini stars and in the four comparison standard stars. The standards have been separated into two groups according to surface gravity to show that their abundances are not dependent on luminosity. but is high compared with the main-sequence Am stars. The [Zr/Fe] ratio appears to become more enhanced with increasing luminosity in the metallic-line stars (Smith 1973c), and again, as in the deficient elements, these anomalous-abundance 8 Delphini stars are similar to the evolved Am stars. Smith (1971) showed that the anomalous [N/Fe] ratios in the metallic-line stars in general do not vary by more than a factor of 2, even though the [Fe/H] ratio in these stars ranges over a factor of 5. That is, in the Am stars the element abundances, [N/H], are strongly coupled to the [Fe/H] abundance. This same effect is present in the five S Delphini stars under discussion here. The well-determined abundances of Ca i, Sc ii, Ti ii, Cr i, Cr n, Mn i, Y n, and Zr n all show less scatter when normalized to the Fe abundance than does the Fe abundance itself. The abundances of [Ca/Fe], [Sc/Fe], and [Zr/Fe] are all similar to those of the evolved Am stars but not those of the main-sequence Am stars. We argued earlier, for statistical reasons, that the slow rotation for these S Delphini stars is intrinsic and, as a + 1.0 + .8 + .6 + .4 ^ + .2 1 ° -.2 -.4 -.6 -.8 -1.0 i—r T I T Tjnr It Ml! ' 1 T S Del Stars Am Stars Evolved Am Stars l-t fcj HH »-t <-> "Ü <" R < in o a > iT » 3 Fig. 8.—A comparison of the abundances normalized to Fe in the anomalous-abundance S Del stars with those in the Am stars (Smith 1971) and those in five evolved Am stars lying from 1 to 2 mag above the main sequence (Smith, private communication). This diagram shows the anomalous-abundance S Del stars to be very similar to the evolved Am stars rather than to the main-sequence Am stars. Note especially the Ca i, Sc n, and Zr n abundances. © American Astronomical Society • Provided by the NASA Astrophysics Data System 662 KURTZ Vol. 32 consequence, they were probably Am stars when they were on the main sequence because slow rotation is thought to be a sufficient condition for metallicism in the Am domain (Abt and Moyd 1973). This, coupled with the present abundance analyses of HR 1706, HR 2255, HR 3265, HR 6561, and HR 7928 and their positions in the H-R diagram, suggests that they are probably evolved metallic-line stars. Abundance analyses of one of the evolved Am stars plotted in Figure 8, 15 Vul, have been previously performed by Miczaika et al. (1956) and by Farra-giana and van't Veer-Menneret (1971), as well as by Smith (private communication). These studies show 15 Vul to have marginal Am characteristics. In particular, Smith finds [Ca/Fe] = -0.09 and [Sc/Fe] = —0.26, along with similar overabundances of [Sr/Fe], [Y/Fe], and [Zr/Fe]. While this star has been reclassified A4 III (marginal Am?) by Cowley et al. (1969), we point out that it has in the past been classified as Am (Slettebak 1949), and that its abundances are indistinguishable from the abundances of the anomalous-abundance 8 Delpbini stars. We have also compared the anomalous-abundance 8 Delphini stars with several other groups of peculiar stars. They do not appear to be similar to the Ap stars. They do not have the large overabundance of Mn associated with the Hg-Mn stars, and Y is enhanced with respect to Fe rather than deficient as in the Sr-Cr-Eu Ap stars (Adelman 1973). Figure 9 compares the anomalous-abundance 8 Delphini stars with the Ba stars (Warner 1965), and shows many similarities between the two classes. The Fe abundance in the Ba stars is enhanced by 0.15 dex, which may not be significant. The only element for which the abundance plotted in Figure 9 is very different between the two groups is Eu, and we feel this difference may not be significant. The Eu abundance is too large in the 8 Delphini stars due to the previously mentioned effect of the neglect of hyperfine splitting, and the Eu abundance in both groups is quite un- certain as only one or two lines were measured for each star. Probably a real difference between the Ba stars and the 8 Delphini stars which is not apparent in Figure 9 is the C abundance. [C/Fe] seems to be normal in the 8 Delphini stars, whereas Warner inferred from the strength of the C2, CH, and CN bands that it is overabundant in the Ba stars by a factor of 2 to 5 times. This overabundance of C, along with the other abundance anomalies in the Ba stars, is thought to be relatively well understood in terms of s-pro-cessed core material being circulated to the surface during a carbon-burning evolutionary stage. Assuming normal evolutionary tracks, the 8 Delphini stars are in a stage of hydrogen burning during which there is no known mechanism for circulating core material to the surface, and during which significant s-processing is probably not possible as 14N(«, />)14C is expected to absorb most of the free neutrons that may be generated. We conclude that the mechanisms which give rise to the similar abundance patterns for the measured elements in the anomalous-abundance 8 Delphini stars and the Ba stars must be different in origin, and the apparent similarity between the two groups is coincidental. The difference in C abundance, the lack of an s-process, core-to-surface circulation mechanism during hydrogen burning, and the Fe enhancement of 0.5 dex in HR 2255 and HR 6561, all support this thesis. The Ba stars show no large overabundances of Fe, and theoretically cannot do so, as the iron-peak elements are the s-process seed nuclei necessary for generating the observed overabundances of the heavier elements. a) HR 3265 The [Y/Fe] abundance anomaly for HR 3265 of + 1.1 dex shown in Figure 6 is significantly larger than the [Y/Fe] anomalies of the other anomalous- ly i.o .8 .6 .4 . .2 0 '-.2 -.4 -.6 -.8 8 Del Stars Ba II Stars u > u e Fig. 9.—Comparison of the abundances normalized to Fe in the anomalous-abundance S Del stars and the Ba u stars (Warner 1965) showing striking similarities between the two groups. See the text for arguments that the similarities are coincidental. © American Astronomical Society • Provided by the NASA Astrophysics Data System No. 4, 1976 METALLICISM AND PULSATION 663 abundance 8 Delphini stars. We consider this unusually large enhancement to be real. It is even apparent in a visual inspection of the spectrum of HR 3265 and has been previously noted by Smith as cited in Baglin et al. (1973). b) HR 7928, 8 Delphini Four previous abundance analyses are available in the literature for 8 Del itself. In Table 7 we compare the present abundances with those derived by Bessell (1969), Breger (1970), Reimers (1969), and Ishikawa (1973), along with the abundances derived from a reanalysis of Ishikawa's data relative to our standards. Our equivalent widths for HR 7928 are determined from two plates, a McDonald Observatory 2.7 m telescope 8 A mm-1 plate and a KPNO 2.1 m telescope 8.9 A mm-1 plate which are in excellent agreement with each other. Nevertheless, the equivalent widths from these two plates are systematically about 20% less than the equivalent widths derived independently by Bessell, Reimers, and Ishikawa, which are basically in agreement with each other. We can offer no explanation for this discrepancy, but note that its effect is mostly canceled out for the abundances normalized to Fe. Reimers's analysis was done with respect to the Sun, using solar gf values for some ions and absolute gf values for others. Our abundances disagree significantly for four of the five ions for which he used absolute gf values, and we suspect that this difference in oscillator strengths is the reason. In addition, the TABLE 7 Comparison of the Abundances from Different Studies of Delta Delphini THIS STUDY BESSELL (1969) [N/Fe] BREGER REIMERS (1970) (1969) ISHIKAWA (1973) ISHIKAWA* C I .11 -.33 -.04 -.13(1) Al I .12 -.16 Si II .42 -.16 -.21 S I .15 Ca I -.08 -.32 -.23 -.13 -.14 .03(9) Sc II -.40 -.42 -.35 -.08 .04 -.40(5) Ti II -.17 -.09 -.26 -.12 .53 -.14(23) V II -.24 -.18 .49 -.42(4) Cr I -.20 .46 - .08 -.05 .29 •03(7) Cr II -.02 .11 -.11 -.03 .19 -.14(11) Mn I .06 -.04 -.13 .09(6) Co I -.04 .03 .44 -.40(1) Ni I -.02 .29 .04 .28(5) Ni II -.01 Zn I .38 .22 .50 •40(2) Sr II .60 .11 .71 .25(1) Y II .41 .59 .99 .15(2) Zr II .32 .54 1.29 .27(3) Ba II .72 .09 .84 •43(1) La II .14 .33 .56 -.14(2) Ce II .06 1.24 -.27(1) Nd II .17 Sm II .12 Eu II .78 .45 1.52 •55(1) Gd II .24 [Fe/H] .00 normal -.25 -.19 -.19 .21 7320 7200 6950 7100 7200 7320 Log g(cgs) 3.25 3.6 3.2 3.8 3.6 3.25 £t(km/sec) 4.5 4.2 4.3 7.0 4.0 4.5 Comparison HR114 n Lep n Lep Sun Sun HR114 Stars HR4825 a CMi HR4825 HR8120 HR8120 HR8272 HR8272 Analysis fine coarse fine fine fine fine * Abundances derived from the data from Ishikawa's study rerun on our system and compared with our standard stars. The numbers in parentheses represent the number of lines used for each ion. © American Astronomical Society • Provided by the NASA Astrophysics Data System 664 KURTZ Vol. 32 discrepancies for Si n, Sr n, and Ba n are at least partially due to the different microturbulent velocities used, as those abundances are derived from a few strong lines lying on the flat portion of the curve of growth. It is less clear why Sc n and Ni i give disparate abundances, and as Reimers's line list is quite different from ours, we did not feel that deriving abundances from the small number of data in common would provide any understanding of the discrepancies. Ishikawa's analysis was also done with respect to the Sun, using the most recent gf values available. While the atmospheric parameters used in this analysis and Ishikawa's are similar and the trend of abundances beyond the iron peak is similar, the actual numerical abundances are quite different. In order to resolve this difference, we have reanalyzed Ishikawa's data relative to our standards, using only the lines from his list which are in common with ours. The results are basically in good agreement with our abundances. The difference in the [Fe/H] ratio is attributable to the difference in the equivalent width scales. We are also in agreement for the five ions in common with Breger's (1970) model-atmosphere reanalysis of Bessell's (1969) data. Breger's lower [Fe/H] ratio is due to the lower effective temperature he used. There are some large differences in some of the derived abundances in the different analyses of HR 7928 listed in Table 7, which makes all of the abundances seem suspect. Nevertheless, we feel that the internal consistency of the abundances of our standards, the general agreement of the present abundances of HR 7928 with those of Breger and Ishikawa's reanalyzed data, and the similarity of the abundances of the five S Delphini stars HR 1706, HR 2255, HR 3265, HR 6561, and HR 7928, all make the run of abundances which we derived in HR 7928 (and, by implication, the other stars analyzed) believable. In particular, in HR 7928 we feel that the [Sc/Fe] deficiency is real, and the marginal deficiencies of [Ti/Fe], rV/Fe], and [Cr/Fe] may also be real. Sr, Y, Zr, and the rare earths are all enhanced relative to Fe. vi. the abundances of HR 1287, HR 2557, HR 3185, and HR 5017 Three of the S Delphini stars analyzed, HR 2557, HR 3185, and HR 5017, and the F2 IV-V S Scuti star analyzed, HR 1287, do not have readily discernible abundance anomalies ([N/Fe] ratios), although the metallicity, [Fe/H], of the S Delphini stars is high. The derived abundances normalized to Fe and the atmospheric parameters adopted are given in Table 8. a) HR 1287, 44 Tauri As has been previously shown in Figure 6, half of the low v sin i large-amplitude S Scuti variables are classified spectroscopically as S Delphini stars. HR 114, one of our standard-abundance stars and one of the standards used by Smith (1971), is the only member of this group which has been shown to be spectroscopically normal by abundance analysis. HR 1287 is an F2 IV-V star (Morgan and Abt 1972) among this TABLE 8 Abundances Normalized to Fe and the Adopted Atmospheric Parameters in the Stars HR 1287, HR 2557, HR 3185, and HR 5017 HR1287 HR2557 HR3185 HR5017 C I - -.28 -.61 -.33 Al I -.24 - -.36 -.32 -.15 Si II .81 -.31 .96 .58 S I - .08 -.77 -.44 Ca I .10 -.16 -.09 -.09 Sc II -.19 .42 -.33 .14 Ti II .07 .03 -.15 .01 V II -.12 .08 -.18 .07 Cr I .13 .02 -.05 -.04 Cr II .10 .11 -.03 -.05 Mn I .07 -.07 .07 .00 Co I . 13 -.01 .04 .14 Ni I -.14 -.20 , .18 .02 Ni II -.23 -.04 -.22 -.03 Zn I - -.63 .22 -.06 Sr II .18 .47 .02 .00 Y II -.18 .19 .04 .09 Zr II -.01 .08 .07 .05 Ba II -.29 -.71 -.29 .01 La II -.32 .36 -.17 .26 Ce II -.29 .09 -.10 .15 Nd II -.44 .59 -.16 -.09 Sm II -.12 .12 -.06 .02 Eu II -.02 .35 .72 .83 Gd II -.10 .38 .09 .16 [Fe/H] 0.20 0.30 0.54 0.80 Teff(K) 7150 7400 7100 7500 log g(cgs) 3.4 3.4 3.25 3.7 5t(km/sec) 5.0 5.5 6.0 5.0 group with v sin / = 20 km s_1 and a pulsational amplitude of 0.07 mag. Our analysis of this star confirms that spectroscopically normal stars can exist in the low v sin i, high-amplitude sector of Figure 6. The Fe abundance lies within the range of Fe abundance for the standard stars. The large apparent overabundance of Si ii is not significant, as this abundance was determined from one partially blended line. There is a trend in the abundances of Ni and the heavier elements for the [N/Fe] ratio to be low, but there is no apparent pattern to these deficiencies and we do not consider the numbers to be meaningfully different from the standards. b) HR 2557 HR 2557 is classified S Delphini by Cowley and Crawford (1971), while Morgan and Abt (1972) classify it A9 III. Crawford has calibrated Mv for Sc! < 0.28 mag for A and early F stars, and an extrapolation of that calibration indicates that HR 2557 lies about 3.4 mag above the main sequence. The effective temperature and surface gravity derived for this star using Breger's calibration (1974a, b) of b — y,fi, and Ci are in good agreement with the excitation and ionization equilibrium for Fe. We have recently discovered this star to be variable, and a detailed analysis of this variability is in progress. The abundances of HR 2557 seem to be marginally abnormal in a manner similar to the previously discussed anomalous-abundance S Delphini stars. The © American Astronomical Society • Provided by the NASA Astrophysics Data System No. 4, 1976 METALLICISM AND PULSATION 665 [Fe/H] abundance may be slightly enhanced, as may be the [N/Fe] ratio for the rare earths. The apparent deficiencies of [Zn/Fe] and [Ba/Fe] are determined from a single line each and thus are suspect. We feel the normality or abnormality of this star is still open to question. c) HR 3185, p Puppis HR 3185, p Pup, is one of the worst examples of stars for which the b — y, (b — y)0, and /J temperature indices disagree. Breger's calibration of b — y, cx yields Tett = 6850 K, log g = 3.5 for this star, while his calibration of /3, cx yields Tetl = 7100, logg = 3.9. We have derived abundances for p Pup using the higher temperature, which is in better agreement with the Fe excitation equilibrium, and find that the Fe ionization equilibrium yields a surface gravity 0.65 dex lower than the photometrically derived gravity. These discrepancies between the predicted effective temperatures and surface gravities from the uvbyfl photometric indices may be due to line blanketing, reddening, or uncalibrated luminosity effects as previously discussed in § II. In Table 9 we compare our derived abundances for p Pup with three other studies of this star. The [Fe/H] TABLE 9 Comparison of the Abundances Normalized to Fe for HR 3185 from Various Studies THIS STUDY GREENS TEIN ET AL. (1948) BESSELL (1969) BREGER (1970) C I -.61 Al I -.32 Si II .96 S I -.77 Ca I -.09 -.32 -.54 -.25 Sc II -.33 -.96 -.44 -.32 Ti II -.15 -.22 -.08 -.19 V II -.18 -.36 Cr I -.05 .07 -.15 -.05 Cr II -.03 -.11 .14 -.08 Mn I .07 -.24 .09 -.13 Co I .04 -.43 Ni I .18 .43 Ni II -.22 .26 Zn I .22 .38 Sr II .02 -.03 Y II .04 .21 Zr II .07 -.30 Ba II -.29 .03 La II -.17 -.21 Ce II -.10 Nd II -.16 Sm II -.06 {rare earths - -.20 Eu II .72 .63 Gd II .09 [Fe/H] 0.54 0.17 0.36 0.42 Te££(K) 7100 - 6800 7000 log g(cgs) 3.25 2.5 3.2 3.3 ^t(km/sec) 6.0 - 6.6 6.0 Comparison HR114 sun n Lep n Lep Stars HR4825 a CM1 HR8120 HR8272 Analysis fine coarse coarse fine abundance is enhanced, but the [N/Fe] ratios of the elements heavier than Fe are all normal, with the possible exception of [Eu/Fe]. Using the curve-of-growth corrections for the effects of hyperfine splitting in Eu AA4130, 4205 (Hartoog, Cowley, and Adelman 1974), we reduce the [Eu/Fe] ratio to about +0.4 dex. Since this abundance is determined from only two lines and the other rare earths do not show similar overabundances, we cannot interpret the Eu abundance as abnormal. Among the lighter elements the C i, Al i, Si ii, and S i abundances are all poorly determined. The Ca i, Sc ii, the well-determined Ti n, and V ii abundances all appear to be slightly deficient when normalized to Fe in all of the abundance analyses. Even though the apparent deficiencies of these elements are small, the agreement among the different analyses indicates that the effect might be real. HR 3185 has slightly enhanced [Fe/H] and possibly some mild deficiencies among the lighter elements when normalized to Fe. d) HR 5017, 20 Canum Venaticorum HR 5017, 20 CVn, has been analyzed previously with respect to the Sun by Dickens et al. (1971) and Ishikawa (1975). We have reanalyzed the data of Ishikawa which are in common with ours, and the results of all these analyses are listed in Table 10. Ishikawa's abundances normalized to Fe look very similar to the anomalous-abundance 8 Delphini stars, but in the comparison of his data run on our system with respect to our standard stars, these anomalies disappear. We consider the abundance ratios, [N/Fe], from this study, from Dickens et al, and from Ishikawa's reanalyzed data to be in basic agreement and conclude that 20 CVn shows no abundance anomalies. The apparent overabundance of Si n is suspect as in the other stars analyzed in this study. If we correct the [Eu/Fe] ratio for the effect of hyperfine splitting, we still find [Eu/Fe] = 0.65. Again, as in HR 3185, this abundance is determined from only two lines and is not accompanied by similar overabundances of the other rare earths, so we do not consider it necessarily anomalous. The [Fe/H] abundance in HR 5017 is enhanced. Dickens et al. (1971) concluded, from their derived abundance, [Fe/H] = 0.44 dex, that the star has metal abundances similar to those of the Hyades. We find [Fe/H] = 0.80 dex, which is a higher metallicity than for the Hyades, and Ishikawa's reanalyzed data give a similar [Fe/H] = 0.67 dex. The equivalent width scales of these studies are not in agreement, however. The abundances in this study are derived from the equivalent widths from two McDonald Observatory 2.1m telescope 8.6 A mm-1 plates which are in excellent agreement with each other, and from one KPNO 2.1 m telescope 8.9 A mm-1 plate which has a slightly larger equivalent width scale. We have adopted the lower scale of the McDonald plates. Our equivalent widths are 23% larger than Ishikawa's which were measured from Okayama Observatory 1.9 m telescope 4 A mm-1 plates. This is typical of © American Astronomical Society • Provided by the NASA Astrophysics Data System 666 KURTZ Vol. 32 TABLE 10 Comparison of the Abundances Normalized to Fe for HR 5017 from Various Studies THIS STUDY DICKENS ET AL. ISHIKAWA ISHIKAWA* (1971) (1975) C I -.33 -.30 -.46(1) Al I -.15 .13 .53(2) Si II .58 .18 S I -.44 -.04 -.40(2) Ca I -.09 -.12 .13 -.13(7) Sc II .14 .23 .24 -.12(7) li II .01 .23 .37 -.09(19) V II .07 .09 .38 Cr I -.04 .11 .21 -.03(1) Cr II -.05 -.02 .09 Mn I .00 -.12 .07 Co I .14 .16 .37 -.49(1) Ni I .02 .01 .07 .08(5) Ni II -.03 Zn I -.06 .35 .11(2) Sr II .00 .12 .57 -.04(2) Y II .09 .28 .19 -.44(1) Zr II .05 -.31 .49 -.30(1) Ba II .01 .31 .70 .14(1) La II .26 -.26 .23 .05(1) Ce II .15 .03 .55 Nd II .09 Srn II .02 Eu II .83 .87 ■30(1) Gd II .16 [Fe/H ] 0.80 0.44 0.33 0.67 Ieff(K) 7500 7520 7875 7500 log g(cgs) 3.7 4.1 3.8 3.7 Ět(km/sec) 5.0 2.0 3.5 4.0 Comparison HR114 sun sun HR114 Stars HR4825 HR8120 HR8272 HR4825 HR8120 HR8272 * Abundances derived from the data from Ishikawa's study rerun on our system and compared with our standard stars. The numbers in parentheses represent the number of lines used for each ion. the difference in equivalent widths usually found between dispersions of 8.6 A mm-1 and 4 A mm-1 (cf. Smith 19736). Our equivalent widths are 43% larger than those of Dickens et al. measured from Mount Wilson 2.5 m telescope 6.8 A mm-1 plates, which is not typical and for which we have no explanation. The difference in the [Fe/H] abundance between this analysis and Ishikawa's reanalyzed data is entirely due to the difference in the respective equivalent width scales. This difference is probably attributable to the differing spectral dispersions used. Since the abundances in this analysis are with respect to standard-star abundances derived from plates of similar dispersion, we consider that the derived metallicity of HR 5017, [Fe/H] = 0.80 dex, is correct to within an estimated internal error of ± 0.2 dex, although some decrease in this ratio is expected due to differential line blanketing between HR 5017 and the four standard stars. The metallicity index, hnix = — 0.035, for this star is in excellent agreement with our [Fe/H] ratio according to the calibration of [Fe/H] with respect to j8 and for Am stars (Rydgren and Smith 1974). vii. duplicity among the bright S delphini stars In Table 11 we list five S Delphini stars which are known to be binary and one which is thought to be. Most of the others have not been tested yet for duplicity, so it is not possible to make a statement about the binary incidence for the group as a whole. Subdividing the class, we note from Table 11 that three of the five anomalous abundance 8 Delphini stars, HR 1706, HR 6561, and HR 7828, are known to be short-period binaries. This is consistent with the interpretation that these stars are evolved metallic-line stars, most of which are binary (Abt 1961; Conti and Barker 1973). We are presently engaged in a program to determine the binary frequency among the bright S Delphini stars. viii. discussion We have suggested that the anomalous abundance 8 Delphini stars, HR 1706, HR 2255, HR 3265, HR 6561, and HR 7928, are evolved metallic-line stars on the basis of their abundances, position in the (fi, Mv) plane, rotational velocity, and binary incidence. Three of these five stars, HR 1706, HR 3265, and HR 7928, are also 8 Scuti pulsators. The other two have not been tested for pulsation. What is the explanation for the anomalous-abundance 8 Delphini stars which seem to be evolved metallic-line stars and 8 Scuti pulsators both? There are several possibilities: (i) the suggestion that the anomalous-abundance S Delphini stars are evolved Am stars may be incorrect, i.e., the abundance anomalies in these 8 Delphini stars may arise from a different mechanism than the abundance anomalies in the Am stars; (ii) each of the pulsating anomalous-abundance 8 Delphini stars may be a binary consisting of an Am star and a 8 Scuti star; (iii) diffusion and pulsation may be able to coexist in a single star under some conditions; or (iv) the diffusion hypothesis may not be the correct explanation for the abundance anomalies of the metallic-line stars. We discuss each of these possibilities below. If the 8 Delphini stars are not evolved Am stars, then there must be two mechanisms for producing TABLE 11 Binaries Among the Delta Delphini Stars HR P (days) Source 1706 3.789 0.004 Harper 1934 2100 2.74050 0.060 Nadeau 1952 4760 Perhaps binary Frost et al. 1929 6561 2.292285 Young 1911 7928 , 40.58 ~i Preston (private communica- tion) 8322 1.022768 0.037 Batten 1961 © American Astronomical Society • Provided by the NASA Astrophysics Data System No. 4, 1976 METALLICISM AND PULSATION 667 very similar abundance anomalies in the 8 Delphini stars and in the metallic-line stars. We have already rejected the s-process mechanism thought to produce the abundance anomalies in the Ba n stars as the source of the anomalous abundances in the 8 Delphini stars. Van den Heuvel (1968a, b) suggested that the Am stars were originally secondary components in binary systems in which the primary evolved and transferred nuclear processed material onto the secondary, which then became Am. This was rejected for the Am stars because they are found in young clusters (Conti 1967; Conti and Strom 1968; Smith 1972a), because Am stars exist in double-line spectroscopic binaries in which the secondary is not a white dwarf, because such a process should not produce the abundance deficiencies of C, Ca, or Sc observed in Am stars, and because the Am stars form a natural continuance of the main-sequence binary frequency. The first objection cannot be applied to the 8 Delphini stars. None of the anomalous-abundance 8 Delphini stars are in clusters, so we can place no age constraint on them in that manner. Preston (private communication) reports that 8 Del itself is a double-line spectroscopic binary (for 4 days out of its 40.58 day period) with nearly identical components. Since the presence of a third close, undetected, evolved component in this system seems unlikely, and since Sc is deficient in 8 Del itself, we reject the mass transfer hypothesis for 8 Del itself and also for the anomalous-abundance 8 Delphini stars as a group. Brancazio and Cameron (1967) suggested that surface nuclear spallation reactions could be responsible for the observed abundance anomalies in the magnetic Ap stars, but Adelman (1973) has argued that this mechanism could be at best only partially responsible for the observed anomalies in the Sr-Cr-Eu Ap stars. Such reactions require large magnetic fields which are present in the cool Ap stars. The magnetic nature of the 8 Delphini stars has not been investigated. We can, however, rule out surface spallation reactions as the source of the anomalous abundances in the S Delphini stars on the basis of the observed abundance pattern. Brancazio and Cameron predict that such reactions should probably enhance Sc relative to Si and enhance Cr and Mn relative to Fe for a bombardment, or should deplete Fe, Co, and Ni and produce an odd-even effect in the run of abundances for pure proton bombardment. None of these abundance relations are observed in the 8 Delphini stars. Havnes and Conti (1971) proposed a magnetic accretion model for the Sr-Cr-Eu Ap stars which, as was later argued by Adelman (1973), is untenable on the basis of the observed abundances in those Ap stars. It does, however, predict a similar overabundance of Si and the iron-peak elements, a greater overabundance of Sr, Y, and Zr, with perhaps Sr and Y being enhanced more than Zr, and an even larger enhancement of the rare earths. These abundance predictions are qualitatively similar to the observed abundance patterns in the anomalous-abundance 8 Delphini stars. We consider magnetic accretion to be unlikely for the 8 Delphini stars because (i) if the 8 Delphini stars are magnetic, then the mechanism operating in the Ap stars should also be present in the 8 Delphini stars, producing Ap-type abundances which are not observed; and (ii) many of the magnetic Ap stars are spectrum variables, whereas none of the 8 Delphini stars is known to be. A Zeeman analysis of the spectra of most of the 8 Delphini stars is not possible due to their relatively high rotational velocities, but could be done on a few of the slowest rotators. Such an analysis would provide a very strong test of the magnetic accretion hypothesis, as the lack of magnetic fields would rule it out and the presence of magnetic fields would show the 8 Delphini stars to be significantly different from the metaliic-line stars which have no measurable fields. The star 32 Vir, a reportedly pulsating classical Am star, has been shown to be a binary in which the primary is a stable Am star and the secondary a normal-abundance 8 Scuti pulsator (Kurtz et al. 1976). A similar explanation for the anomalous-abundance 8 Delphini stars is very attractive since at least three of them are binary and it eliminates the need to explain how a star can be metallic-lined and also pulsate. Unfortunately, this hypothesis is probably not tenable for the 8 Delphini stars. In none of our 8-10 A mm-1 spectra of these 8 Delphini stars is there any indication of line doubling. As has been mentioned, 8 Del itself is an SB2 system in which both components are 8 Scuti pulsators. The [Sc/Fe] deficiency in 8 Del itself is difficult to explain with this model, and, finally, one of the properties of such a system should be different radial velocity curves for the metallic lines and the Ca n K line, as the Am star would dominate the metallic-line spectrum while both components would contribute to the Ca n K line, as is the case for 32 Vir. We tested for this effect in 13 8 Delphini stars on 24 8-10 A mm-1 plates by measuring radial velocities from those plates using the KPNO Grant measuring machine. In no case was the K line velocity significantly different from the metal-line velocity. We therefore reject the Am- 8 Scuti star binary hypothesis as a model for the anomalous-abundance 8 Delphini stars. We are left with the choice that either (i) diffusion is a much stronger phenomenon than previously thought and can exist in a pulsating star, or (ii) diffusion is not the correct theoretical explanation of the Am phenomenon. We prefer the first alternative because of the success of the diffusion hypothesis in explaining the observed properties of the metallic-line stars. The second is tantamount to the null hypothesis for the Am stars. Where do the high-metallicity, but non-anomalous-abundance, 8 Delphini stars, HR 2557, HR 3185, and HR 5017, fit into the above scheme? HR 2557 could be interpreted as having transition abundances between the anomalous-abundance 8 Delphini stars and normal stars, and HR 3185 and HR 2557 are the most luminous, and therefore probably the most evolved, of the 8 Delphini stars. It is tempting, therefore, to postulate that these three stars are representative of the last phase of the transition of metallic-line © American Astronomical Society • Provided by the NASA Astrophysics Data System 668 KURTZ Vol. 32 stars into normal stars. They are, however, inter-pretable as having abundances within the cosmic scatter of metallicity for normal stars. This is certainly true for HR 2557 and HR 3185. While for HR 5017 our metallicity is too high to be in the range of normal stars, other investigators (Dickens et al. 1971) found a Hyades-type metallicity for this star. The unusual abundances for these three stars, coupled with the similarity of HR 3185 and HR 5017 to the anomalous abundance 8 Delphini stars HR 1706, HR 3265, and HR 7928 in their rotational velocity and pulsational amplitude and in their position in the (jS, Mv) plane, suggest that all of these stars may have a common origin. ix. suggestions concerning the relationship between metallicism and pulsation We have in the previous section stated our preference for retaining the diffusion hypothesis as a working model for the Am stars, and we have suggested that the pulsating anomalous-abundance 8 Delphini stars are evolved Am stars in which pulsation and metalli-cism coexist. We propose the following diffusion model to explain the relationship between metallicism and pulsation. Following the suggestion of Watson (1971) and Smith (1971, 1973a), we propose that diffusion occurs in the radiative zone between the H i, He i, and He n ionization zones in the Am stars. This qualitatively predicts the run of abundances observed in the Am and Fm stars, especially the [Ca/Fe] and [Sc/Fe] deficiencies. Following Breger (1972), Baglin (1972), and Vauclair et al. (1974), we suggest that helium is sufficiently depleted from the He n ionization zone to inhibit pulsation, thus accounting for the observed exclusion between the classical Am stars and the 8 Scuti pulsators. This also serves to reduce the convective overshoot from the He n ionization zone into the overlying radiative zone, thus mitigating the objections of Latour et al. (1975) to diffusion occurring in this zone. We further suggest that the helium depletion in the He n ionization zone is not sufficient to eliminate convection in that zone. This provides a convective barrier between the upper and lower radiative zones so that the Am anomalies can arise quickly from diffusion in the upper radiative zone and then remain essentially time-independent during the main-sequence lifetime of the Am star as required by the observations (Smith 1972a, 1973a). Enough helium remains in the He n ionization zone that pulsational instabilities can grow in an Am star as it evolves into the giant region as is required by the evidence presented in this paper that some Am stars evolve into 8 Scuti pulsators. The pulsating anomalous-abundance 8 Delphini stars are examples of these evolved Am stars in which either (i) diffusion occurs below the He n ionization zone where the pulsational amplitude becomes small due to the increasing density, or (ii) mixing across the upper radiative zone is a slow enough process that these 8 Delphini stars represent evolved Am stars in which diffusion no longer occurs but for which mixing has not yet eliminated the apparent surface anomalies. Pulsation and metallicism may coexist in other border regions of the Am domain. Abt (1975) has shown that while rotation, temperature, and age are sufficient to determine if a star will be metallic-lined or not, they are insufficient to determine the strength of metallicism in a given metallic-line star. Some other factor or factors are involved, and we have hypothesized that pulsation may be one of them. We are presently in the process of testing all of the marginal Am stars for pulsation, and have found two so far, HR 4594 and HR 8210, which do pulsate. In summary, we propose that (i) classical Am stars do not pulsate, (ii) metallicism and pulsation can coexist among the subgiant and giant A and F stars as in the anomalous-abundance 8 Delphini stars, and (iii) pulsation and metallicism may coexist among the marginal (Am:) metallic-line stars. I would like to express my sincere thanks to Dr. Michel Breger, who suggested and supervised this work, for his continued guidance and help. I am equally grateful to Dr. Myron Smith for many enlightening discussions and for very generously provided unpublished abundances of HR 6561 and five evolved Am stars. Thanks are due to Drs. Myron Smith, Michel Breger, Leonard Kuhi, Deane Petersen, and Mr. Frank Fekel for generously providing some of the spectra used in this work. It is a pleasure to acknowledge KPNO for the use of the PDS micro-densitometer, Grant measuring engine, KPNO radial velocity reduction program, and the spectrophoto-metric reduction programs, SPECT1 and SPECT2. I also thank Dr. Robert Kurucz for making WIDTH5 available. This work was submitted to the University of Texas in partial fulfillment of the requirements of the degree of Doctor of Philosophy. APPENDIX A In this appendix we clarify the meaning of the various subclassifications of the metallic-line and related stars. Classical Am stars are stars which are classified Am according to the MK classification criteria defined by Roman, Morgan, and Eggen (1948). This usually means that the K-line type and the metal-line type differ by five or more subclasses. The hydrogen line types which are intermediate between the K-line types and metal-line types for these stars range from A4 through Fl, and are consistent with the derived temperatures. The classical Am (or Fm) stars are therefore metallic-line stars with pronounced line-strength anomalies. One should note, however, that misclassification may occur, as in the case of 15 Vul which Slettebak (1949) called Am by the Roman, Morgan, and Eggen criteria (hence we would call it classical Am), but which has been reclassi* © American Astronomical Society • Provided by the NASA Astrophysics Data System No. 4, 1976 METALLICISM AND PULSATION 669 fied A4 III (marginal Am?) by Cowley et al. (1969). The abundance analyses referenced in this paper support the latter classification. In addition, there are Am stars with pronounced abundance anomalies determined from curve-of-growth analyses which are not classified as classical Am stars, namely the early Am stars. Early or hot Am stars are stars earlier than A4 which Conti (1965) pointed out have pronounced Am anomalies as evidenced by the Sc n A4246/Sr n A4215 Ľne ratio. So far as is known, these stars are phenomenologically the same as the classical Am stars, but are not classified as classical Am because, at the surface temperature of the early A stars, the H fines are at their broad maximum, the K line is on the flat portion of the curve of growth, and the metal line strengths are weakening due to increased ionization, making the MK criteria insensitive to abundance anomalies. Marginal or mild Am stars are Am stars in which the difference between the K-line type and the metal-line type is less than the five subclasses necessary for classification as a classical Am star. There is a selection effect in the stars classified as marginal Am (Am:) by Cowley et al. (1969) in that the marginal Am stars are systematically hotter than the classical Am stars according to the b — y or temperature indicators of the uvbyfi photometric system. There is therefore probably a large overlap in the hot Am and the marginal Am classifications. A photometric analysis of the marginal Am stars is presently in progress, and a complete discussion of this classification will be treated in a future publication. Delta Delphini stars are stars with spectra similar to S Del; that is, stars with subgiant and giant luminosity types and metal-line spectra similar to the Am stars. The S Delphini classification is a spectroscopic classification only. We have shown in this paper that physically many of the S Delphini stars are probably evolved Am stars. FIVp-FIIp stars are early F subgiant and giant stars which are classified as peculiar by Morgan and Abt (1972) in their definition of the MKA system because the K-line type is earlier than the metal-line type. Malaroda (1973, 1975) uses the MKA criteria to classify stars as S Delphini, and Anne Cowley (private communication) agrees that for classification purposes these MKA F IVp-F lip stars are 8 Delphini stars. Morgan and Abt call 8 Del itself FO IVp. Care should be taken not to confuse the MKA Fp stars with the very probably physically unrelated magnetic peculiar stars such as y Eql or 49 Cam which are classified FOp (Cowley et al. 1969). Delta Scuti stars are Population I short-period pulsating A and F stars within three magnitudes of the main sequence. The 8 Scuti classification implies pulsational variability; it does not imply anything spectroscopic about a star. See Baglin et al. (1973) for a complete discussion. APPENDIX B EQUIVALENT WIDTH DATA FOR THE ABUNDANCE ANALYSES PRESENTED IN THIS STUDY Table 12 is a fisting of the stars and plate material used for the differential fine abundance analyses presented in this paper. The plates were traced using the KPNO PDS microdensitometer and converted to intensity using the KPNO spectrophotometric reduction programs SPECT1 and SPECT2 on the University of Texas CDC 6600 computer. Photographic density-to-intensity calibrations were used at 4100 and 4630 Á. All plates were IIa-0 emulsion with projected slit widths of 20 fan and widening of 0.4-0.8 mm. Equivalent widths were measured treating all lines as triangles. The equivalent width data derived are presented in Table 13 along with the excitation potentials and oscillator strengths used which are from the lists of Corliss and Warner (1964), Corliss and TABLE 12 Plate Material Used for the Abundance Analyses Presented in This Study Star Dispersion (A mm-1) Telescope Name HR Plate Number Observatory (m) Observer 44 Tau 1287 EC6245 10 Lick 3.1 Kuhi 14 Aur 1706 Ce20609a 10 Mt. Wilson 2.5 Petersen B8925 8.6 McDonald 2.1 Kurtz 6 Mon..... 2255 B8507 8.6 McDonald 2.1 Smith B8924 8.6 McDonald 2.1 Kurtz 2557 B8503 8.6 McDonald 2.1 Smith P Pup...... 3185 B8774 8.6 McDonald 2.1 Fekel 3265 EC6249 10 Lick 3.1 Kuhi 20CVn.... 5017 D2011 8.9 KPNO 2.1 Smith B8929 8.6 McDonald 2.1 Kurtz B8930 8.6 McDonald 2.1 Kurtz S Del 7928 D2843a 8.9 KPNO 2.1 Breger N597 8.0 McDonald 2.7 Kurtz 28 And____ 114 N621 8.0 McDonald 2.7 Kurtz y Vir 4825 D2847 8.9 KPNO 2.1 Breger 8120 N613 8.0 McDonald 2.7 Kurtz 8272 N610 8.0 McDonald 2.7 Kurtz © American Astronomical Society • Provided by the NASA Astrophysics Data System 1976ApJS...32..651K @ > re s* p 3 > as o o 3 re SL o re re" H3 c EL re Q. re > in > > o B* re" a — to re 3 TABLE 13 Measured Equivalent Widths for the Program Stars (mA) o x(eV) log gf HR114 HR4825 HR8120 HR8272 C I 4771. 72 7 46 -1. 70 59 32 51 119 4775. 85 7 46 -2. 20 - 17 18 25 4932. 00 7. 65 -1. 92 62 - 37 22 Al I 3944. 01 0. 00 -0. 62 195 143 131 202 3961. 52 0. 01 -0. 32 193 160 115 192 Si II 4128. 05 9. 79 0. 22 144 35 127 - 4130. 88 9. 80 0. 77 57 42 106 108 S I 4694. 13 6. .50 -1. 39 10.6 19 8.8 23 4695. 45 6 .50 -1. .54 - - 13.4 30 4696. 25 6. .50 -1. 76 12.4 10.5 - 19 Ca I 4226. 72 0 .00 -0 .55 295 - 161 316 4283. 01 1 .88 -0 .39 147 69 37 114 4302, 53 1 .89 0 .30 172 162 95 - 4318. 65 1 .89 -0 .15 - - - 60 4425. 43 1 .87 -0 .33 103 74 34 84 4435. 69 1 .88 -0 .69 65 53 31 79 4455. 89 1 .88 -0 . 72 95 82 33 83 4578. 56 2 . 51 -0. .82 26 48 21 23 4585. 87 2 . 52 -0. .31 67 62 32 54 Sc II 4246. 83 0 .31 0 .09 232 - - 255 4294. 77 0 .60 -1 .27 39 25 38 38 4305. 70 0 . 59 -1. .33 164 - - - 4314. 08 0 62 -0. .10 191 117 231 207 4320. 76 0 .60 -0. .22 232 138 239 231 4325. 01 0 59 -0 .37 147 97 134 - 4374. 45 0 .62 -0. 45 - 100 - 171 4400. 36 0. .60 -0 80 156 82 135 4415. 56 0. .59 -0 94 - - - - 4431. 37 0. .60 -.2 02 25 - 21 17 Ti II 3913. 46 1. 12 -0 .24 271 179 192 238 4012. 37 0. .57 -1. .58 178 100 178 186 4028. 33 1. 89 -0. 65 123 70 111 140 4056. 21 0. .61 -2. .46 - 24 30 50 4163. 63 2. 58 0, 45 - 98 - 181 4294. 10 1. .08 -0. .90 196 131 188 196 4300. 05 1. 18 -0. .46 269 209 253 A (A) x(eV) log gf HR114 HR482S HR8120 HR8272 4 301 .93 1 .16 -1. 11 136 89 1 5 6 4312. .86 1 .18 -1. 06 159 111 149 192 4367. . 06 2 . 69 -0. 39 121 - 11 4 .1 31 4386.86 2 .60 -0. 46 77 48 - 78 4390. .98 1. . 23 -2. 03 82 - - 76 4394. .06 1 .22 -1. 47 100 47 98 116 4395. .03 1. .08 -0. 50 235 179 198 241 4395, .85 1. .24 1. 53 67 45 78 85 4399. .76 1 .24 -1. 06 165 82 140 164 4411. .08 3. .09 -0. 07 57 31 85 74 4411. .94 1 .22 -2. 11 28 19 4 6 - 4417. ,72 1. . 16 -1. 18 181 92 154 179 4418. . 34 1 .24 -1. 67 106 61 82 116 4421. ,95 2 .06 -1. 14 - 19 60 77 4443. .80 1. .08 -0. 74 Ill 118 191 87 4450. . 49 1 .08 -1. 41 92 111 1 29 170 4464. .46 1 .17 -1. 66 102 68 118 169 4468. . 49 1. .73 -0. 65 218 105 175 24(1 4488 . 32 3 .12 0. 01 76 56 92 99 4493 .53 1. .08 -1. 92 24 12.6 - 4501 .27 1. .12 -0. 79 232 152 20 1 240 4529. .46 1. .57 -1. 52 - 78 108 99 4533 .97 1 .24 -0. 64 - 207 261 277 4544 .01 1 .24 -2. 08 26 31 47 - 4545. .14 1. . 13 -1. 61 46 22 64 6 7 4563. .76 1 . 22 -0. 86 - 107 20 7 2 34 4568. 31 1. . 22 -1. 93 29 27 - - 4571. .97 1 . 57 -0. 34 260 162 212 249 4589 .96 1 . 24 -1. 61 129 113 129 128 4 7 79.99 2. . 04 -1. 12 - 20 1 14 1.0 2 4805 .11 2 .05 - 0 . 74 - - 168 153 V II 400 2 .94 1 .43 -1. 28 38 11.9 - - 4005. ,71 1 .82 -0. 22 - - - - 4008. . 1 7 1 . 7 9 -1, 61 17 8.7 17 22 4023. 39 1. .80 -0. 35 61 37 80 62 4036. 78 1 .48 -1. 42 12.7 17 32 2 2 4039 5 7 1 .82 -1. 60 9.8 22 14 - 4183. .43 2'. .05 -0. 77 16 34 31 5 5 Cr I 3919.16 1. .03 0.14 94 51 4 5 - 4254.35 0 00 -0. 27 149 - 108 190 4274.80 0. .00 -0. 39 127 118 95 1 49 4371. 28 1, .00 -0. 90 33 50 - - 4511. 90 3. .07 0. 37 20 9.1 - 25 4591. 39 0. .96 -1. 18 15 - 10.6 - 4616. .13 0. .98 -0. 95 57 24 48 - 4646.17 1. .03 -0. 49 73 64 2 3 - 4651.28 0. .98 -0. 97 33 22 9.9 1 7 4652. 16 1. 00 -0. 78 34 44 - 28 4664. 20 3. 11 0. 37 - 11.5 - - 4718, 45 3, 18 0. 82 35 8.1 9.9 36 1976ApJS...32..651K TABLE 13—Continued © > f?" s > © ss © 3 SL O -i © a" sr t/i O sr r?" Ö p: re \(X) eV) log : gf HR114 HR4825 HR8120 HR8272 X(A) x(eV) log gf HR114 HR4S25 HR8120 HR8272 Cr II Fe I 4209. 35 3. 83 -1. 75 - 18 23 - 4047'. 32 2, .28 -1. ,84 7.8 - - 13.0 4252. 16 3. 86 -1. 85 31 17 57 54 4049. 33 2. .59 -1. ,35 21 22 - - 4261. 92 3. 86 -1. 21 43 47 106 109 4059 . 72 3, ,54 -0. ,52 7.1 11.9 - 39 4269 . 30 3. 85 -2. 06 36 - - - 4062. 44 2, ,84 0. 05 - 56 44 - 4275. 58 3. 86 -1. 33 34 25 80 94 4063. 54 1. .56 0. ,43 27 245 175 313 4284. 21 3. 86 -1. 85 38 21 57 59 4065. 39 3. .43 -0. , 70 - 18 - - 455S . 02 4. 07 -1. 44 58 35 70 85 4067. 98 3. .21 0. ,29 74 80 36 96 4558. 66 4. 07 -0. 45 156 90 196 182 4070. 76 3. .24 0. ,01 64 49 - - 4588. 22 4. 07 -0. 65 144 80 - 157 4071. 74 1. .61 0. ,40 237 217 - 217 4592. 09 4. 07 -1. 37 74 80 91 99 4072. 52 3. ,43 -0. ,43 - 28 26 - 4616. 64 4. 07 -1. 51 57 24 48 81 4073. 76 3. . 14 -0 . 14 64 33 25 51 4618. 82 4. 07 -0. 98 144 96 108 133 4074. 79 3, ,05 -0. ,14 43 47 - 69 4634. 11 4. 07 -1. 19 111 - 105 112 4076. 73 3, , 21 0. ,24 149 170 - 150 4848 . 24 3. 85 -1. 13 74 59 - 115 4079. 84 2, ,86 -0. ,52 40 - 11.0 - 4876 . 41 3. 86 -1. 94 76 42 " 4084. 4091. 49 56 3. 2, ,33 .83 0. -1. 13 .28 49 81 Mn I 4107. 49 2. .83 0. .06 70 34 23 60 4030 . 77 0. 00 -0. 48 - - 138 240 4112 . 97 4. ,18 -0. ,02 37 52 - - 4033. 07 0. 00 -0. 64 182 206 129 175 4114. 44 2, ,83 -0. .47 37 - 21 - 4034. 49 0. 00 -0. 88 125 90 61 119 4120. 21 2. .99 -0. ,43 26 31 - 24 4035. 73 2. 11 0. 37 108 106 99 121 4123. 74 2, ,61 -1. ,13 18 - - - 4041 . 36 2. 11 0. 93 65 108 - 113 4126. 19 3, ,33 -0. 35 29 64 - 42 4048. 76 2. 13 0 . 25 69 95 - 133 4132. 06 1. .61 -0. ,16 166 172 140 186 4055. 54 2. 13 0. 47 31 SO 25 - 4132. 90 2. .84 -0. ,02 63 - - 80 4082. 94 2. 17 0. 25 - 17 - 65 4133. 86 3. .37 -0. 48 42 - 21 35 4083. 62 2. 13 0. 24 25 83 28 99 4134. 68 2, ,83 0 . .18 105 98 65 97 4502. 22 2. 91 0. 18 18 22 - 19 4136. 51 3. .37 -0. ,82 8.8 12.6 8.0 - 4754. 04 2. 27 0. 13 34 79 - - 4137. 00 3. ,14 0. 12 87 70 - 68 4783. 42 2. 29 0. 11 " 56 82 4139. 4140 . 93 44 0. 3. .99 .42 -2. -1. 86 ,11 14.7 12.6 14.7 8.5 19 Fe I 4141. 86 3. .02 -1. 04 18 - - 27 3815. 84 1. 48 0. 60 279 302 - - 4147. 67 1. ,48 -1. 47 89 100 22 25 3865. 82 1. 01 -0. 56 225 129 161 - 4153. 90 3. ,40 0. 33 127 151 - 117 3871. 80 2. 95 -0. 15 95 70 41 - 4156. 80 2. .83 0 . 13 108 172 - - 3872 . 50 0 . 99 -0 . 54 - - 199 - 4157. 78 3. , 42 0. 17 76 69 - 72 3895. 65 0 . 11 -1. 47 - 147 183 - 4174. 92 0. 91 -2. 34 26 70 11.0 20 3902. 94 1. 56 0. 12 273 157 - 244 4175. 64 2. ,84 0. 10 - 61 33 - 3920 . 26 0 . 12 -1. 49 177 124 117 190 4176. 57 3. 37 0. 04 45 79 37 64 3922. 91 0. 05 -1. 41 194 112 125 181 4181. 75 2. 83 0 . 46 184 180 - 181 3927. 92 0. 11 -1. 29 240 152 160 198 4182. 38 3. 02 -0. 37 18 73 - - 3955. 35 3. 28 -0. 41 - 50 37 - 4184. 89 2. 83 -0. 05 - 66 45 - 3983. 96 2. 73 0. 06 99 102 69 4187 . 04 2 . 45 0. 17 119 128 70 139 3998. 05 2. 69 -0. 04 123 96 43 - 4187. 80 2. 42 0. 13 142 147 113 153 4005. 24 1. 56 -0. 07 - 251 - - 4191. 43 2. 47 0. 06 182 147 - 145 4007. 27 2. 76 - 0. 45 37 47 24 - 4202. 03 1. 48 -0. 25 218 175 158 228 4017 . 15 3. 05 -0. 17 90 94 - 79 4203. 98 2. 84 -0. 21 101 89 51 75 4021. 87 2. 76 0. 12 114 69 - 85 4207. 13 2. 83 -0. 69 44 59 - - 4029 . 64 3. 26 -0. 42 39 44 34 48 4210. 35 2. 48 -0. 19 106 - 71 - 4040 . 65 3. 30 -0. 30 41 49 28 - 4213. 65 2. 84 -0. 55 33 31 17 51 4043. 90 2. 73 -0 . 56 84 70 29 94 4216. 29 0. 00 -2. 98 64 48 - 47 4044. 61 2. 83 -0 . 17 51 68 19 - 4217. 55 3. 43 0 . 12 76 78 47 94 4045. 81 1. 48 0. 68 377 249 333 4219. 4222. 4225. 36 21 46 3. 2. 3. 57 45 42 0. -0. 0 . 79 35 13 125 85 100 101 65 79 63 65 120 112 1976ApJS...32..651K TABLE 13—Continued © > 2_ f?" v S > v. © © 3 SL O •1 o a" t/i O sr r?" Ö '< re OS -J to XCÄ) x(eV) log gf HR114 HR4825 HR8120 HR8272 Fe I 4227. 43 3. 33 0. 90 164- - 140 206 4228. 72 3. 37 -1. 65 - - - - 4235. 94 2. 42 0. 31 192 ■- 129 181 4238. 81 3. 40 0. 47 - 55 80 114 4240 . 37 3 . 55 -0. 59 26 - 18 - 4245. 35 2. 86 -0. 44 54 47 34 89 4246. 09 3. .64 -0. 42 26 22 17 4247. 43 3. 37 0. 45 90 82 - - 4248. 22 3. .07 -0. 53 23 29 12.3 39 4250 . 12 2 , .47 0 . 25 - 102 89 160 4250. 79 1. .56 -0. 28 186 125 114 195 4260. 47 2. .40 0. 63 255 247 - 257 4264. 21 3, .37 -0. 78 20 13.4 26 4266. 96 2. .73 -0. 87 - 25 9.6 39 4267. 98 3. . 11 -0. 34 42 33 18 32 4268. 74 3. .30 -0. 63 16 40 11.9 - 4271. 15 2. .45 0. 25 132 - 98 - 4271. 76 1. .48 0. 20 207 219 186 242 4276. 68 3. .88 -0. 75 - 12.7 - - 4282. 41 2. .18 -0. 16 - - 98 124 4285. 44 3. .24 -0. 42 34 42 - 42 4291. 46 1. .56 -1. 99 21 22 78 - 4298. 04 3. ,05 -0 . 56 24 43 - 32 4325. 76 1. .61 0. 36 147 214 178 19S 4327. 92 3. .30 -0. 90 - - - - 4352. 74 2, ,21 -0. 56 77 42 - - 4375. 93 0, , 00 -2. 59 91 - 49 92 ' 4376. 78 3. ,02 -1. 27 11.9 8.7 9.3 11.7 4382. 78 3. ,57 -0. 16 - - 14.7 - 4383. 56 1, ,48 0. 51 - 299 - - 4387. 87 3, ,07 -0. 62 22 50 - - 4388. 41 3. .60 0. 02 68 74 - - 4392. 58 3. .88 -0. 96 - 9.2 - - 4404. 75 1, .56 0. 25 - 199 156 223 4408. 41 2, .20 -0. 95 58 91 38 70 4415. 12 1. .61 -0 . 13 - - - - 4422. 57 2. .84 -0. 22 - 56 - 97 4427. 31 0. .05 -2. 51 120 - - 111 4430. 61 2 .22 -1. 02 56 83 - 90 4432. 57 3. .57 -0. 82 - 21 8.8 - 4433. 22 3 .65 -0. 14 55 - - 36 4438. 35 3 .69 -0. 99 17 8.4 - - 4442. 34 2 .20 -0 . 50 138 - 65 - 4443. 19 2 .86 -0. 22 111 - - 87 4447. 72 2 .22 -0. 58 90 66 67 88 4466. 58 2 .83 0. 18 132 - 78 103 4469. 38 3 .65 0. 19 115 78 88 - 4476. 02 2 . 84 0. 14 127 102 77 137 4479. 61 3 .69 -0. 70 - 34 17 - 4480. 14 3 .05 -1. 01 - - 14.8 - 4484. 23 3 .60 0. 08 53 59 37 46 4485. 67 3 .68 -0. 40 30 28 - 43 4490. 08 3 .02 -0. 74 28 32 - ■ x(eV) log gf HR114 4494.57 4495.97 4517.53 4525.14 4531.63 4547.85 4587.13 4602.00 4602.94 4607.67 4611.28 4625.05 4630.11 4632.92 4635.84 4637.51 4638.01 4643.46 4647.43 4668.14 4673.17 4678.86 4690.17 4691.42 4705.46 4707.28 4710.28 4733.60 4735.85 4736.78 4745.81 4772.82 4122.63 4128.73 4178.85 4233.16 4273.31 4296.56 4303.16 4351.76 4385.38 4416.81 4472.92 4489.18 4491.40 4508.28 4515.33 4520.24 4522.63 1.61 1. 48 2. 81 2.84 -0.35 -0.99 -1.11 0.03 -1.20 -0.10 -0.91 -2.50 -1.46 -0.66 -0.13 -0.63 -1.83 -2 . 31 -1 .46 -0.60 -0 . 35 -0. 59 -0.47 -0 .30 -0.53 0 .05 -0. 79 -0.54 -1.27 -0.23 -0.74 -2.38 -0.37 -0.02 -0.55 -1.05 -2.73 -2. 76 -2.00 -1.43 -2.27 -2.36 -2.00 -1.76 -2.09 -1.76 -1.91 -1.87 -1.51 24 13.4 22 32 49 64 32 60 57 58 35 25 89 57 145 132 177 155 115 66 132 178 166 211 8S 18 25 77 26 26 26 28 60 42 25 9. 38 10.9 42 12.6 71 35 37 21 95 14.7 23 137 201 70 82 109 35 93 87 82 96 11.2 7.4 12.0 18 19 24 112 72 178 106 133 149 218 190 135 75 156 156 178 168 149 207 106 12.7 53 21 69 42 23 51 30 32 16 65 18 32 81 104 68 210 125 139 176 161 156 87 124 140 191 191 174 233 1976ApJS...32..651K TABLE 13—Continued © > IT S > © ES c 3 s O 13 -s o < a re Q- sr re > > Vi o TS sr M. S" Ü ta re ON X(eV) log gf HR114 HR4825 HR8120 HR8272 A (A) x(eV) log gf HR114 HR4825 HR8120 HR8272 Fe II Zr II 4541. 52 2. 85 -2. 29 - .51 - 132 4149. .22 0 .80 -0 .13 , 118 127 98 133 4555. 89 2. 83 -1. 79 246 - 215 238 4150. ,97 0 .80 -1 .02 - 10 .1 18 38 4576. 33 2. 84 -2. 22 - 50 127 125 4156. .24 0 . 71 -0 .85 79 50 49 79 4582. 83 2. 84 -2. 44 97 72 102 - 4208. ,99 0 . 71 -0. .54 41 - 51 64 4583. 82 2. 81 -1. 25 242 168 245 244 4211. ,88 0 .52 -1 .21 12, ,7 10 . 5 43 27 4620. 51 2. 83 -2. 63 78 52 93 56 4317. 32 0. . 71 -1 .48 24 6 .9 8.9 . 4629. 33 2. 81 -1. 78 166 . 70 158 187 4635. 35 5. 95 -1. 43 - - - 38 Ba II 4656. 97 2. 89 -2. 53 135 66 - 149 4554. 03 0 ,00 0 .17 220 - 135 204 4666. 75 2 . 83 -2. 64 83 70 - 10 2 4934. 10 0 .00 -0 ,14 173 183 124 4670. 17 2. 58 -2. 87 - 41 110 89 4731. 44 2. 89 -2. 29 142 79 - 92 La II 4923. 92 2. 88 -0. 93 253 - - - 3988. 51 0 ,40 -0. .26 11. 5 - 16 11 4086. 72 0. .00 -0, .60 16 26 24 - Co I 4123. 23 0. .32 -0, .40 22 25 24 20 4020. 90 0. 43 -1. 58 7.1 11.1 - 7.1 4263. 59 1. .95 0, .03 (4. 8) - 10.2 17 4058. 60 2. 00 -0. 67 26 54 13.8 51 4322. 51 0. .17 -1. .62 - 4121. 32 0. 92 -0. 03 39 32 24 33 4333. 76 0, .17 -0, .60 52 1 4 ,1 - - 4662. 51 0. .00 -2, .04 (4. 5) - 14.1 - Ni I 4748. 73 0 ,92 -1. .20 5, 4 9.7 13.1 3858. 30 0, 42 -0. 62 166 132 87 - 4401. 55 3. 18 0. 83 - 98 39 87 Ce II 4606. 99 3. .47 0. 77 - 55 30 39 3882. 45 0, 32 -0. .06 . 56 133 4606. 23 3. 58 0. 30 - 16 23 9.9 4120. 83 0, ,32 -0. ,74 (4. 21 - - 4648. 65 3. 47 0. 78 66 41 13.1 42 4137. 63 0, ,04 0. ,09 27 29 14.8 8.2 4686. 22 3. . 58 -0. 39 21 12.2 - - 4142. 40 0 , ,22 -0. ,14 31 45 13.8 - 4714. 42 3, ,37 0 . 84 88 45 47 76 4193. 09 0. ,74 -0. ,10 - (4. 0) 4.5 - 4715. 78 3 .53 0. 76 34 - 13.9 20 4202. 94 0. ,56 -0. 34 7. 4 5.7 10.5 4418. 76 0. ,38 0. 03 25 13. 6 17 - Ni II 4486. 91 0. 29 -0. 62 10. 9 6. 4 5.0 5.8 4015. 50 4. .03 -1. 25 63 57 34 54 4067. 05 4. .03 -0. 59 132 157 125 176 Nd II 4244. 80 4. .03 -2. 03 - 34 - 29 4061. 09 0. 47 0. 03 17. 8 (5. 0) 18 4362. 10 4. .03 -1. 43 - 14.7 37 - 4462. 98 0. 56 -0. 84 22 25 - 110 Zn I Sm II 4722. 22 4. .01 0. 69 29 35 5.4 31 4424. 34 0. 48 -0. 42 . 14. 0 9.2 16 4810. 53 4. .06 0. 86 14.9 27 19 50 4467. 34 0. 66 -0. 39 7. 4 6. 7 - Sr II Eu II 4077. 71 0, ,00 -0. 78 283 276 216 309 4129. 73 0. 00 -0. 31 20 16. 8 12.0 13 4215. 52 0 00 -0. 99 248 241 214 268 4205. 05 0. 00 -0. 08 61 46 - Y II Gd II 3950. 35 0. .10 -0. 71 111 - - 80 4251. 73 0. 38 -0. 39 - (4.7 10.2 12.5 3982. 59 0. . 13 -0 . 79 65 45 82 87 4177. 54 0. .41 -0. 24 186 196 196 210 4309. 62 0. .18 -0. 98 106 163 96 - 4358. 73 0, .10 -1. 61 24 54 37 23 4374. 94 0. .41 -0. 14 - - - 171 4398. 02 0 .13 -1. 25 49 - - 53 197 6ApJS...32. . 65 TABLE 13—Continued © > f?" s > o o 3 SL O •1 o < a" Z o TS sr r?" Ö p: re on -j 4^ ACS) HR1287 HKL706 HR2255 HR2557 HR3185 HR3265 HR5017 HR7928 A (A) HR1287 HR1706 HR2255 HR2557 HR3185 HR3265 HR5017 HR7928 C I Ti II 4771.72 - 31 106 64 - 122 143 76 4301, ,93 192 190 240 255 276 236 253 121 4775.85 _ 23 49 29 37 34 44 41 4312. ,86 210 147 210 197 256 159 234 120 4932.00 79 _ _ 24 _ - 30 4367, ,66 184 115 184 194 260 162 202 98 4386. ,86 113 70 129 111 168 111 171 60 Al I 4390, ,98 146 74 153 131 216 146 198 70 3944.01 204 184 221 212 266 247 247 191 4394, ,06 144 96 136 123 175 99 178 57 3961.52 178 131 220 148 224 162 199 171 4395, ,03 266 219 260 274 326 240 302 186 4395. ,85 124 54 130 100 143 196 175 43 Si II 4399, ,76 193 140 204 217 252 64 232 117 4128.05 135 120 162 196 279 - 204 . 89 4411, ,08 107 57 94 73 143 - 149 41 4130.88 75 95 164 98 259 206 230 107 4411, .94 82 29 56 48 101 26 110 13.4 4417, .72 192 150 180 167 250 175 216 118 S I 4418, .34 131 49 99 - 153 99 158 46 4694.13 _ 21 49 45 65 18 37 43 4421, .95 100 53 84 - 138 53 142 20 4695.45 16.6 26 29 23 10.1 37 14.9 4443, ,80 227 187 216 - 292 - 248 173 4696.25 15.3 17 35 11.0 5.2 18 20 4450, .49 183 140 212 169 - 222 249 113 4464, ,46 168 106 161 160 266 166 212 80 Ca I 4468, .49 227 179 240 - 299 222 245 163 4226.72 363 289 319 338 396 369 364 268 4488, ,32 139 63 116 105 180 88 166 50 4283.01 148 97 128 - 167 126 158 93 4493, .53 62 22 46 61 81 22 84 93 4302.53 195 166 162 210 240 - 195 123 4501, .27 228 176 260 190 296 234 269 164 4318.65 151 _ - 116 - 120 - - 4529 .46 165 55 149 119 196 103 195 53 4425.43 146 78 120 99 162 85 158 79 4533, .97 290 252 310 297 352 298 329 - 4435.69 139 68 151 91 210 - 173 95 4544, .01 65 26 72 107 82 31 96 23 4455.89 132 51 120 - 168 122 146 79 4545, ,14 84 34 93 96 - 53 127 37 4578.56 57 25 52 - 65 39 81 40 4563. ,76 239 156 233 217 260 205 261 155 4585.87 124 37 98 87 118 65 146 43 4568. ,31 - 18 37 47 - 10.1 83 - 4571, ,97 260 196 298 260 325 302 319 198 Sc II 4589, ,96 142 95 161 - 192 92 187 77 4246.83 197 222 249 - 280 249 274 155 4779, .99 - 75 88 120 - - 154 74 4294.77 84 45 74 62 95 44 128 IS 4805, .11 - 102 162 - 220 - 180 112 4305.70 - - - 220 - - - - 4314.08 212 196 280 254 325 278 303 163 V II 4320.76 - 219 281 290 333 261 316 141 4002, ,94 72 141 54 134 - 155 29 4325.01 180 112 192 197 233 - 232 94 4005, ,71 116 - 162 - - - 197 80 4374.45 178 - 201 302 205 - 219 91 4008. .17 - 21 42 40 - - 98 12.3 4400.36 155 106 196 153 177 179 237 81 4023, ,39 110 81 131 96 150 118 165 55 4415.56 153 - 154 - 169 - 188 70 4036, ,78 38 36 57 33 66 36 97 11.5 4431.37 26 14.8 41 62 50 10.1 79 8.6 4039, .57 18 27 26 55 52 13.8 47 8.4 4183, ,43 71 39 77 - 104 111 128 18 Ti II 3913.46 266 238 278 319 362 325 314 217 Cr I 4012.37 221 188 278 212 290 298 307 170 3919 ,16 158 64 138 - 186 130 170 74 4028.33 146 149 172 185 179 156 196 89 4254, .35 207 164 216 - 277 196 248 140 4056.21 62 52 66 - - 30 142 12.8 4274 .80 204 130 192 - 276 203 221 137 4163.63 188 145 199 - 257 176 232 111 4371, .28 115 49 92 - 99 121 104 44 4294.10 231 198 234 203 299 196 244 175 4511, ,90 31 13.2 24 - 47 15 36 17 4300.05 264 225 285 250 372 298 311 181 4591, ,39 48 14.2 46 - 50 26 66 7.7 4616, ,13 64 54 52 97 92 92 19 4646, ,17 - 36 101 92 170 - 129 44 4651, .28 14.7 40 52 70 18 59 24 4652, .16 - 26 47 41 94 34 77 32 4664, ,20 - 14.8 12.8 - - - - - 4718 ,45 - 18 28 46 69 - 64 20 1976ApJS...32..651K TABLE 13—Continued > 2 2, r.' to s > © = c 3 SL o -s c Q." ft — a- sr re > o -a sr o" bs Ö B9 i-f-P J/5 2 ON -J HR1287 HR1706 HR2255 HR2557 HR3185 HR326S HR5017 HR7928 HR1287 HR1706 HR2255 HR2557 HR3185 HR3265 HRS017 HR7928 Cr II Fe 1 4209.35 - 14.9 23 - - 44 - - 4047.32 10.1 22 31 42 7.6 29 15 4252.16 63 43 69 - 165 85 140 36 4049.33 21 49 82 47 73 23 4261.92 126 95 149 124 211 142 193 73 4059.72 26 58 27 110 47 86 29 4269.30 68 33 108 92 134 110 157 35 4062.44 92 167 86 185 162 184 89 4275.58 123 60 116 103 176 99 161 56 4063.54 275 253 339 368 405 368 286 267 4284.21 99 51 81 98 153 99 143 34 4065.39 55 24 80 66 34 79 21 4555.02 109 56 88 125 146 85 141 74 4067.98 136 97 124 186 123 173 85 4558.66 194 143 210 - 240 204 239 119 4070.76 110 73 118 _ 150 114 151 62 4588.22 156 124 176 149 205 137 199 67 4071.74 272 210 263 244 262 267 276 192 4592.09 114 78 112 - 141 141 157 19 4072.52 67 28 59 72 85 _ 111 19 4616.64 101 35 99 101 136 78 133 51 4073.76 101 59 125 129 _ 129 65 4618.82 156 78 144 176 220 126 192 93 4074.79 115 48 113 71 148 no 147 61 4634.11 134 77 122 144 185 92 161 83 4076.73 152 188 214 240 257 139 4848.24 - - 110 - 159 - 138 77 4079.84 84 28 67 90 _ - 105 55 487C.41 - - 167 - 165 - 173 ■ 4084.49 120 49 96 126 66 123 55 »1 I 4091.56 29 - 19 15 29 6.1 29 - 299 218 4107.49 97 38 114 80 141 85 108 70 4030.77 - - - - - 289 4112.97 35 74 77 115 43 107 31 4033.07 230 177 219 262 290 270 177 4114.44 72 29 72 37 76 49 96 36 4034.49 176 133 173 140 231 187 204 124 4120.21 80 23 67 80 103 62 103 35 4035.73 171 145 178 137 237 170 228 94 4123.74 96 32 106 89 121 72 147 54 4041.36 158 113 155 141 232 162 203 104 4126.19 102 34 95 141 70 119 38 4048.76 134 118 179 137 227 205 227 115 4132.06 256 160 263 275 314 319 294 188 4055.54 98 75 96 91 168 83 130 51 4132.90 117 77 101 188 148 149 80 4082.94 87 13.1 73 60 96 S3 97 - 4133.86 84 59 120 163 107 142 54 4083.62 147 48 120 109 177 114 176 51 4134.68 166 126 162 279 162 209 113 4502.22 36 8.9 33 39 63 16 38 89 4136.51 34 37 52 56 19 4754.04 - 27 87 - 171 75 102 51 4137.00 112 59 109 79 197 . 156 73 4783.42 - 43 101 73 107 107 60 4139.93 10. S 23 22 70 20 57 - Fe 1 4140.44 33 12.0 35 23 65 22 59 14.6 313 303 4141.86 43 13.5 61 - . - 86 19 3815.84 - - - - - 4147.67 128 62 114 129 169 122 151 71 3865.82 - - 254 - - 240 - - 4153.90 164 110 133 _ 249 _ 217 99 3871.80 122 - 212 - 188 214 208 - 4156.80 169 97 143 - 224 - 220 124 3872.50 - - - - - 245 - 4157.78 130 80 94 93 186 117 154 82 3895.65 202 - 268 - 293 232 175 4174.92 95 147 95 78 162 62 112 44 3902.94 232 221 340 296 - 331 256 208 4175.64 139 80 116 75 205 141 159 91 3920.26 190 165 225 188 228 231 217 162 4176.57 130 61 108 107 192 116 151 74 3922.91 224 167 263 228 283 278 265 189 4181.75 158 198 193 328 249 254 165 3927.92 220 177 275 225 264 312 245 216 4182.38 85 43 86 33 162 90 127 49 3955.35 108 37 105 52 132 73 153 54 4184.89 111 73 105 71 187 114 143 93 3983.96 148 86 167 - 228 179 222 103 4187.04 178 149 176 281 267 219 138 3998.05 158 88 163 138 264 187 208 99 4187.80 230 141 221 181 314 249 251 147 4005.24 116 261 - - - - 275 210 4191.43 206 140 229 258 262 141 4007.27 83 45 123 - 124 - 125 34 4202.03 240 217 279 250 379 285 310 192 4017.15 122 87 128 182 159 146 171 - 4203.98 157 74 127 91 224 137 201 89 4021.87 148 95 155 143 239 - 196 215 121 4207.13 78 39 72 66 174 114 149 41 4029.64 101 93 115 - 137 142 140 77 4210.35 163 100 168 225 170 185 112 4040.65 98 51 126 112 205 - 174 - 4213.65 64 44 76 51 146 82 no 54 4043.90 127 62 123 - 196 117 170 74 4216.29 110 29 118 48 205 157 80 4044.61 108 59 96 - 151 99 147 58 4217.55 no 59 139 41 192 130 148 80 4045.81 431 314 420 436 537 428 462 360 4219.36 164 98 143 107 234 161 196 99 4222.21 143 104 161 153 249 201 179 104 4225.46 113 129 156 160 - 214 225 103 1976ApJS...32..651K TABLE 13—Continued © > f?" S > VI © © 3 SL O •1 o a" t/i O sr r?" Ö '< Vi re -a HR1287 HR1706 HR2255 HR2557 HR3185 HR3265 HR5017 HR7928 XCA) HR1287 HR1706 HR2255 HR2557 HR3185 HR3265 HR5017 HR7928 Fe I 4227.43 209 165 215 217 273 279 279 155 Fe I 4494.57 174 93 177 165 243 175 183 118 4228.72 - 21 - 19 25 - 34 4495.97 - 10.0 42 40 - 14.2 48 10.9 4235.94 228 169 244 186 325 305 256 187 4517.53 43 16 39 56 72 22 57 14.9 4238.81 159 100 156 109 257 - 201 101 4525.14 128 65 138 113 242 150 175 99 4240.37 69 29 49 41 - 31 110 25 4531.63 29 - 20 - - - 34 - 4245.35 127 49 103 - 200 114 162 61 4547.85 69 23 61 97 107 53 97 42 4246.09 67 24 72 131 92 114 36 4587.13 45 11.4 25 27 42 - 52 10.9 4247.43 171 96 151 101 184 154 189 111 4602.00 - 14.1 33 - - 13.8 46 17 4248.22 74 31 71 79 - - 106 43 4602.94 122 43 89 51 153 75 130 63 4250.12 179 126 183 - 268 - 209 124 4607.67 - 16.2 71 50 84 75 92 45 4250.79 202 167 212 - 269 - 239 160 4611.28 118 34 85 73 160 72 143 69 4260.47 256 214 266 367 279 268 196 4625.05 87 22 77 66 116 59 102 43 4264.21 44 21 34 32 92 - 74 23 4630.11 - 18 38 15 54 18 56 35 4266.96 50 30 51 59 115 42 73 23 4632.92 64 21 53 - - 26 88 32 4267.98 94 33 68 55 130 56 96 37 4635.84 35 24 30 - 55 26 27 38 4268.74 71 26 58 - - - 99 36 4637.51 71 37 46 71 115 - 94 ' 4271.15 199 150 202 - 250 - 225 143 4638.01 68 29 72 60 118 - 94 62 4271.76 268 202 285 - 354 - 277 208 4643.46 44 40 57 - - 34 74 36 4276.68 31 16 34 46 - 13.6 57 20 4647.43 - - 91 85 180 78 136 64 4282.41 165 117 198 145 253 214 222 130 4668.14 - - 84 105 167 - 106 63 4285.44 87 28 80 - 143 88 114 45 4673.17 - 29 72 83 94 36 85 - 4291.46 53 17 43 67 103 21 70 - 4678.86 - 40 78 116 - 85 113 80 4298.04 81 27 65 81 - - 86 32 4690.17 21 20 43 46 - 47 19 4325.76 292 206 315 280 359 - 308 228 4691.42 - - 76 75 167 - 137 54 4327.92 25 - 30 30 - 16 59 10.9 4705.46 - - 24 35 - - 21 15 4352.74 137 47 112 72 164 67 147 73 4707.28 - 54 83 80 184 - 115 65 4375.93 123 54 128 116 185 82 146 86 4710.28 - 31 60 51 100 - 84 37 14376.78 22 23 23 31 53 11.6 38 5.6 4733.60 - - 42 71 - - 68 28 4382.78 - 39 44 48 82 53 80 32 4735.85 - 11.3 35 - 80 - 53 4383.56 282 220 292 288 357 - 283 247 4736.78 - - 108 77 - - 140 80 62 4387.87 64 43 59 55 122 107 110 28 4745.81 - 12.1 60 50 - - 58 4388.41 103 40 92 82 159 68 146 52 4772.82 - 30 50 - - 61 45 4392.58 16 - 14.7 41 47 - 57 5.6 4404.75 255 189 262 257 335 - 280 218 Fe II 95 4408.41 - 59 105 - 204 141 169 68 4122.63 163 99 168 160 197 141 208 4415.12 225 - 230 335 - 261 201 4128.73 135 57 89 111 123 89 149 - 4422.57 131 58 145 201 146 167 96 4178.85 185 160 216 225 273 263 256 152 4427.31 160 89 150 - 217 156 196 107 4233.16 266 260 259 358 372 - 304 225 4430.61 129 60 89 105 192 130 164 68 4273.31 123 96 169 140 209 164 183 82 4432.57 19 17 16.8 - 47 - 52 13.0 4296.56 184 128 220 189 266 222 243 140 4433.22 78 34 75 83 121 82 111 46 4303.16 186 147 189 202 268 236 217 145 4438.35 17 11.1 25 33 52 10.1 54 9.5 4351.76 - - - 298 - 309 - - 4442.34 147 75 81 127 193 - 163 93 4385.38 188 134 214 - 294 170 240 141 4443.19 112 89 135 - 199 - 164 115 4416.81 179 130 182 - 231 168 204 127 4447.72 137 60 111 97 184 127 152 .98 4472.92 157 47 133 - 216 147 172 60 4466.58 148 104 155 117 207 162 186 119 4489.18 172 102 170 145 251 146 218 106 4469.38 - 84 141 232 145 202 88 4491.40 170 127 166 142 233 149 190 119 4476.02 177 85 145 121 236 154 186 123 4508.28 190 117 215 167 278 190 233 160 4479.61 37 18 49 48 - - 83 27 4515.33 192 115 197 151 264 179 220 150 4480.14 - 13.6 45 - - - 80 17 4520.24 182 128 182 176 248 170 215 143 4484.23 - 33 87 - 131 114 115 63 4522.63 228 166 254 210 305 303 301 196 4485.67 64 14.4 71 45 - 48 76 28 4490.08 69 28 54 - - 62 82 33 5a SS S3 8 "Kg^S.K ST S3 S iiiiii u umm mum m u m i HESS S 8 S 3 SSSSgSS.SSgS. SKR SK<°3 Sg S3RSSSRK£KS3S! SSK SKSSSSSiK SSSS SK RSS3RSS iiiiiiiiii nam jiii ii ii mm 677 American Astronomical Society • Provided by the NASA Astrophysics Data Syste 5 10 100 500 o Equivalent Width (mA) KPNO Fig. 10.—Comparison of the equivalent widths measured by the author for S Del from a McDonald 2.7 m telescope 8.0 A mm-1 plate and a KPNO 2.1 m telescope 8.9 A mm-1 plate. 500l i i i i i o< J. 100 !3 -D — o i: L±J 10 -I-1—1—1—I I I I I ... .yy*'' j_i_i 11111 10 100 o Equivalent Width(mA) KPNO 500 Fig. 11.—Comparison of our equivalent widths for HR 114 from a McDonald 2.7 m telescope 8.0 A mm-1 plate with the equivalent widths for HR 114 of Smith (19726) from a KPNO 2.1 m telescope 8.9 A mm-1 plate. © American Astronomical Society • Provided by the NASA Astrophysics Data System METALLICISM AND PULSATION 679 500 0< E 100 «_ o c Z SS a. 5 * ..XX •• x**x? / •% • • X x • . * **- • X» x * • Mt.Stromlo x Mt. Wilson 100 o Equivalent Width(mA) 500 Fig. 12.—Comparison of our equivalent widths for 8 Del with those of Bessell (1969) and Reimers (1969). Bessell and Reimers agree with each other, whereas our equivalent widths are smaller than theirs. Note, however, that our equivalent widths from two plates of 8 Del are in good internal agreement, as shown in Fig. 10. Bozman (1962), Lambert and Warner (1968a, b, c, d), and Warner (1967). The oscillator strengths used are not the most recent but rather were chosen to be the same as those used by Smith (1971, 1973a) for most of the line list in order to facilitate intercomparison of the S Delphini star abundances and Smith's Am and Fm star abundances. Because all abundances were derived differentially, errors in oscillator strength values cancel out to first approximation. A critical discussion of more recent oscillator strengths useful in the analysis of A and F stars is given by Ishikawa (1973, 1975). We have compared the equivalent-width data used in this analysis with other published equivalent-width data for the same stars and find on the average no systematic shift. The mean scatter of our equivalent widths compared with others is +0.09 dex. In Figures 10,11, and 12 we have plotted some representative examples comparing our data from different telescopes and with the data of other investigators. Figure 10 compares our equivalent widths for 8 Del from a plate taken with the McDonald 2.7 m telescope at 8 A mm-1 and a plate taken with the KPNO 2.1 m telescope at 8.9 A mm-1. Figure 11 compares our equivalent widths for HR 114 with those of Smith (1971). Both of these figures are typical of most of the equivalent-width comparisons made both internally and with the published equivalent widths of other investigators. Figure 12 represents a worst case. In it, equivalent widths of 8 Del from a McDonald 2.1 m telescope 8.6 A mm-1 plate are compared with those measured by Bessell (1969) at 6.8 A mm-1 at Mount Stromlo and Reimers (1969) at 10 A mm-1 at Mount Wilson. There is a systematic shift of our equivalent widths compared with theirs of about 20yo or 0.08 dex. No explanation of this shift is offered. Bessell and Reimers's data are in good agreement but the author's data shown in Figure 12 are also compared with another plate of 8 Del in Figure 10, with good agreement there, also. We have used a heterogeneous group of plates from different telescopes and observatories, but find that the equivalent-width scales derived from plates of comparable dispersion are usually in good agreement with those of other investigators, regardless of the equipment used to obtain the plates. The rms scatter in the compared equivalent widths is + 0.09 dex. In some cases our equivalent widths differ from those in the literature by as much as 20%. Although such a shift in the equivalent-width scale is exceptional in this analysis, it must be kept in mind in any interpretation of the data presented herein that such a shift may be present. Special care should be taken especially in the case of stars for which only one plate was measured. REFERENCES Abt, H. A. 1961, Ap. J. Suppl., 6, 37. -. 1965, Ap. J. Suppl., 11, 429. -. 1975, Ap. J., 195, 405. Abt, H. A., and Moyd, K. I. 1973, Ap. J., 182, 809. Adelman, S. J. 1973, Ap. J., 183, 95. Allen, C. W. 1963, Astrophysteal Quantities (2d ed.; London: Athlone). Baglin, A. 1972, Astr. Ap., 19, 45. Baglin, A., Breger, M., Chevalier, C, Hauck, B., Le Contel, J. M., Sareyan, J. P., and Valtier, J. C. 1973, Astr. Ap., 23, 221. Batten, A. H. 1961, Pub. D.A.O., 11, 395. Bessell, M. S. 1969, Ap. J. Suppl., 18, 167. Bidelman, W. P. 1965, Vistas in Astronomy, 8, 53. © American Astronomical Society • Provided by the NASA Astrophysics Data System 680 KURTZ Brancazio, P. J., and Cameron, A. G. W. 1967, Canadian J. Phys., 45, 3297. Breger, M. 1968, Pub. A.S.P., 80, 578. -. 1970, Ap. J., 162, 597. -. 1972, Ap. J., 176, 373. -. 1974a, Ap. J., 192, 75. -. 19746, Dudley Obs. Rept., 9, 31. Conti, P. S. 1965, Ap. J., 142, 1594. -. 1967, in The Magnetic and Related Stars, ed. R. Cameron (Baltimore: Mono), p. 321. -. 1970, Pub. A.S.P., 82, 781. Conti, P. S., and Barker, P. K. 1973, Ap. J, 186, 185. Conti, P. S., and Strom, S. E. 1968, Ap. J., 152, 483. Corliss, C. H., and Bozman, W. R. 1962, N.B.S. Monog. No. 53. Corliss, C. H., and Warner, B. 1964, Ap. J. Suppl., 8, 395. Cowley, A. P. 1968, Pub. A.S.P., 80, 453. -. 1973, Pub. A.S.P., 85, 314. Cowley, A. P., and Cowley, C. R. 1964, Pub. A.S.P., 76, 119. -. 1965, Pub. A.S.P., 77, 184. Cowley, A. P., Cowley, C. R., Jaschek, M., and Jaschek, C. 1969, A.J., 74, 375. Cowley, A. P., and Crawford, D. L. 1971,Pub. A.S.P., 83, 296. Cowley, A. P., and Fraguelli, D. 1974, Pub. A.S.P., 86, 70. Crawford, D. L., and Barnes, J. V. 1970, A.J., 75, 978. Crawford, D. L., Barnes, J. V., Faure, B. Q., Golson, J. C, and Perry, C. L. 1966, A.J., 71, 709. Danziger, I. J., and Dickens, R. J. 1967, Ap. J., 149, 55. Danziger, I. J., and Faber, S. M. 1972, Astr. Ap., 18, 428. Dickens, R. J., French, V. A., Owst, P. W., Penny, A. J., and Powell, A. L. T. 1971, M.N.R.A.S., 153, 1. Edmonds, F. N., Schlutter, H., and Wells, D. C, III. 1967, Mem. R.A.S., 71, 271. Eggen, O. J. 1956, Pub. A.S.P., 68, 238. Faraggiana, R., and van't Veer-Menneret, C. 1971, Astr. Ap., 12 258 Frost, E. B., Barrett, S. B., and Struve, O. 1929, Pub. Yerkes Obs., 7, 1. Harper, W. E. 1934, Pub. Dominion Ap. Obs., 6, 207. Hartoog, M. R., Cowley, C. R., and Adelman, S. J. 1974, Ap.J., 187, 551. Hauck, B. 1971, Astr. Ap., 11, 79. Havnes, O., and Conti, P. S. 1971, Astr. Ap., 14,1. Hoffleit, D. 1964, Catalogue of Bright Stars (3d ed.; New Haven: Yale University Press). Ishikawa, M. 1973, Pub. Astr. Soc. Japan, 25, 111. -. 1975, Pub. Astr. Soc. Japan, 27, 1. Kukarkin, B. V., Efremov, Yu. I., and Kholopov, P. N. 1958, General Catalog of Variable Stars (Moscow: IAU). Kurtz, D. W., Breger, M., Evans, S. W., and Sandmann, W. H. 1976, Ap. J., 207, 181. Kurucz, R. L. 1970, S.A.O. Spec. Rept., No. 309. Lambert, D. L., and Warner, B. 1968a, M.N.R.A.S., 138, 183 --. 19686, M.N.R.A.S., 138, 213. -. 1968c, M.N.R.A.S., 138, 229. -. 1968a", M.N.R.A.S., 140, 197. Latour, J., Spiegel, E. A., Toomre, J., and Zahn, J.-P. 1975, Bull. A.A.S. (abstract), 7, 526. Lindemann, E., and Hauck, B. 1973, Astr. Ap. Suppl., 11,119. Malaroda, S. 1973, Pub. A.S.P., 85, 328. -. 1975, A.J., 80, 637. Michaud, G. 1970, Ap. J., 160, 641. Miczaika, E. R., Franklin, F. A., Deutsch, A. J., and Green-stein, J. L. 1956, Ap. J., 124, 134. Milton, R. L., and Conti, P. S. 1968, Ap. J., 154, 1147. Morgan, W. W., and Abt, H. A. 1972, A.J., 77, 35. Nadeau, P. H. 1952, Pub. David Dunlap Obs., 1, 537. Reimers, D. 1969, Astr. Ap., 3, 94. Roman, N. E., Morgan, W. W., and Eggen, O. J. 1948, Ap. J., 108, 107. Rydgren, A. E., and Smith, M. A. 1974, Ap. J., 193, 125. Slettebak, A. 1949, Ap. J., 109, 547. Smith, M. A. 1971, Astr. Ap., 11, 325. -. 1972a, Ap. J., 175, 765. -. 19726, Astr. Ap. Suppl., 5, 81. -. 1973a, Ap. J. Suppl., 25, 277. -. 19736, Ap. J., 182, 159. -. 1973c, invited paper presented for the Joint Session of Comm. 29 and 36 of the IAU, Sydney, Australia. -. 1976, Ap. J., 203, 603. van den Heuvel, E. P. J. 1968a, Bull. Astr. Inst. Netherlands, 19, 309. -. 19686, Bull. Astr. Inst. Netherlands, 19, 326. Vauclair, G., Vauclair, S., and Pamjatnikh, A. 1974, Astr. Ap., 31, 63. Warner, B. 1965, M.N.R.A.S., 129, 263. -. 1967, Mem. R.A.S., 70 165. Watson, W. D. 1970, Ap. J. (Letters), 162, L45. -. 1971, Astr. Ap., 13, 263. Young, R. K. 1911, Lick Obs. Bull., 6, 160. Donald W. Kurtz: Department of Astronomy, San Diego State University, San Diego, CA 92182 © American Astronomical Society • Provided by the NASA Astrophysics Data System