The Zeta Aurigae stars

The Zeta Aurigae Stars* K. O. WRIGHT Dominion Astrophysical Observatory, Victoria, B. C. SUMmARy The elements of the orbits of the ~ Aurigae stars are discussed and the masses and dimensions derived. I t is found t h a t the mass ratio, ~ l / ~ s , is less than 1.5 for the two systems, ~ Aurigae and 31 Cygni, where the data are well determined, and masses of approximately 8 (~) for the giant K star and 6 (~) for the main-sequence B star are found. The spectral types and luminosities of the component stars are derived after a discussion of the photoelectric and spectrographic observations. These data are then used to obtain the dimensions of the stars, on the assumption that the effective temperatures and bolometric corrections are known. The dimensions are also calculated from the eclipse data, where it is assumed that the eclipses of ~ Aurigae and 31 Cygni are almost central. The latter data give diameters 135 (~) for these K stars and 5 (~) for the B stars. The diameters calculated from the luminosities are about 50 percent larger for the K stars, but are of the same order for the B stars; the differences are probably within the uncertainties of the data. The absolute magnitudes of ~ Aurigae and 31 Cygni derived from the eclipse data are between the values given by Hodge and Wallerstein and by Wilson in their calibrations of the emission width of the K line. Observations of the Ca II K line, made within a few months of eclipse, are reviewed for the systems Aurigae, 32 Cygni and 31 Cygni, The intensities of these lines at a given phase differ from one eclipse to another, and at ingress and at egress for the same eclipse. The observations are interpreted in terms of atmospheres which vary in extent from time to time. Several satellite lines can sometimes be seen on the same spectrogram when high dispersion is used. The results confirm detailed studies of the inner chromospheres which require the presence of condensations to explain the observations.

T h e s t u d y of t h e ~ A u r i g a e stars e x t e n d s over a b o u t 40 years, b u t i t has b e e n intensified r e c e n t l y because i t is n o w r e a l i z e d t h a t a n a l y s i s of t h e i r s p e c t r a will give us v a l u a b l e i n f o r m a t i o n concerning t h e a t m o s p h e r e s of s u p e r g i a n t l a t e - t y p e s t a r s t h a t c a n n o t a t p r e s e n t be o b t a i n e d in a n y o t h e r way. These systems, of which ~ A u r i g a e is t h e p r o t o t y p e , consist of a g i a n t l a t e - t y p e s t a r w h i c h , a t r e g u l a r i n t e r v a l s , eclipses a m u c h s m a l l e r earlyt y p e star. One i m p o r t a n t o b s e r v a t i o n a l r e s u l t is t h a t effects of t h e a t m o s p h e r e of t h e g i a n t s t a r on t h e r a d i a t i o n f r o m t h e B - t y p e s t a r c a n be d e t e c t e d for some t i m e o u t s i d e of eclipse. T h e effects c a n b e s t be o b s e r v e d s p e c t r o s c o p i c a l l y b u t , for a c o m p l e t e s t u d y of t h e p h e n o m e n a , p h o t o m e t r i c o b s e r v a t i o n s are also needed. Commission 42 of t h e I n t e r n a t i o n a l A s t r o n o m i c a l U n i o n , d e v o t e d to studies of p h o t o m e t r i c d o u b l e stars, h a s b e e n a c t i v e in p r o m o t i n g c o - o p e r a t i v e observing p r o g r a m m e s in o r d e r t o o b t a i n a d e q u a t e coverage of t h e s e eclipses a n d full c r e d i t s h o u l d be given t o t h e m a n y observers who h a v e o f t e n g i v e n u p t i m e on t h e i r r e g u l a r p r o g r a m m e s in o r d e r t h a t t h e s e s t a r s m a y be o b s e r v e d a t critical phases. * Contribution~ #ore the Dominion Astrophysical Observatory, No. 116. Published by permission of the Deputy Minister, Department of Energy, Mines and Resources, Ottawa, Canada. 147

148

The Zeta Aurigae Stars

The stars usually included in this class are ~ Aurigae, 32 Cygni, 31 Cygni, VV Cephei and, possibly, e Aurigae. Other somewhat similar systems, AR Pavonis, GG Cassiopeiae, K U Cygni, RZ Ophiuchi, V777 Sagittarfi, V381 and V383 Scorpii, BL Telescopii and Boss 1970, have not been studied in any detail because they are much fainter and therefore spectroscopic observations cannot be obtained so readily (see Fracastoro, 1956; PayneGaposchkin, 1960; and Wood, 1963). However, when it proves desirable to check the results obtained for the bright stars, the others can be observed with the large telescopes now, or soon to be, in operation. Time is often a factor in obtaining these observations since, although the periods are long, the B star crosses behind the atmosphere of the cool star quite rapidly and since details concerning the atmosphere of the latter are to be studied, as many observations as possible should be obtained with high dispersion. The technique of single-trail spectra used so successfully by Preston and the Lick observers may be very useful for these stars. Among this group of stars, ~ Aurigae and 31 Cygni are probably the best examples since the spectra of the component stars seem to be quite normal and the eclipses, very probably, are nearly central. The orbit of 32 Cygni is quite inclined and therefore the eclipses, though total, are not far from grazing. The VV Cephei system consists of a supergiant M star and, presumably, a Be star with an emission envelope. Since the spectrum of the M star probably varies, as do those of other similar stars, and since the envelope of the B star complicates the analysis of the observations, an interpretation of the atmosphere of the giant star is much more difficult than for the other stars. However, it is probably worthwhile to make the attempt since the M star is cooler and the atmosphere is more extensive than those of the other systems. The system of e Aurigae is still a puzzle although many explanations for the spectroscopic observations have been given; it will not be discussed in as great detail at this time since few new observations are available. The present contribution will be devoted to a discussion of the orbits and dimensions of these systems, and to observations of the Call K line obtained up to two stellar diameters from the limb of the cool star. Little or no detailed consideration will be given to the inner atmospheres of these stars, derived from the "chromospheric" spectra seen very close to eclipses, as an enhancement of low-excitation lines of other atoms, because Wilson (1960) discussed the observations available at that time quite adequately. A few new observations have been published (e.g. Wright, 1959, for 31 Cygni) but the principal conclusions remain the same and further refinements in the theory are required. The photometric observations will not be examined exhaustively except in the study of the periods, the times of contact required to determine the dimensions of the stars, and the spectral types of the components derived from the colours. Comprehensive bibliographies of recent observations are given in reports of Commission 42 in the I . A . U . Transactions and Reports on Astronomy. The necessary data and the discussion will be given for each star separately, and the results will be summarized at the end of the paper.

THE ORBITS AND DIMENSIONS

(i) Zeta Aurigae The spectroscopic orbit was first studied by Harper (1924) who derived a period of 972 days. Harper himself noted that one plate showed the many-lined spectrum of a K-type star rather than the usual "washed-out" appearance characteristic of the composite spectrum but could not convince himself that an eclipse was taking place because of the

K. O. WRIGHT

149

low probability of such an occurrence and because the date of the observation did not agree within the probable error with the computed time of eclipse. In 1937 Harper revised his orbit (see Lee and Wright, 1960) combining his later observations with Mount Wilson and Lick data, after applying small systematic corrections to the latter. A new orbit has recently been obtained from these same data by means of a computer programme, using Wood's (1951) value of the period, 972.162 days; the velocity curve and the observations fitted to this orbit are shown in Fig. 1, and the elements are given in Table 1. A new KM/SEC./ /

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FIG. 1. Radial velocity curve for ~ Aurigae. value for the mass ratio has been obtained from" studies of the secondary spectrum. Victoria observations made near maximum separation of the components in 1961 combined with Lee and Wright's (1960) data give a semi-amplitude, K2, for the secondary spectrum of 31.4 kin/see and corresponding masses of 7.6 ± 1.2 q) and 5.9 ± 1.5 (D respectively for the primary and secondary stars. Popper (1961} found comparable masses of 8.3 ± 1.5 Q and 5.6 ~ 1.1 Q. Although Lee and Wright's method of obtaining velocities of the secondary star is probably capable of greater accuracy, Popper's values are based on more observations. Therefore mean values, ~ g ~ 8.0 ~ 1.2 Q and ~ B = 5.8 ± 1.1 q) have been adopted for the masses of the ~ Aurigae system. Recent photometric observations of later eclipses suggest that the period may not be precisely constant. Tanabe and Nakamura (1957) derived a mean period of 972.141 days based on eclipses observed from 1934 to 1956; they also suggested that the length of totality has increased about 1 day in that time. Hardorp, Herczeg and Seholz (1966) have adopted a period of 972.153 days based on 1934 to 1963 observations and again there is an indication that the duration of total eclipse is increasing gradually. However, these small differences in the period do not affect the orbital elements appreciably.

1.0

Semi-major axis, k m

Semi-amplitude, km/see

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4585.74 4- 4.46 5478.4 547O.O -[-12.86 4- 0.26 24.58 431.4 43.OO~ 3.83 x

1147.8 0.3014-0.013 218.2 -4-3.0

32 Cygni

7169.73 7660.4 7685.7 --7.73 --12.3 13.98 20.8 7.09 10.55

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3784.3 0.222-4- 0.008 201.1 -4- 2.4

31 Cygni

Orbital elements of the Zeta Aurigae stars

972.162 0.406 4- 0.011 336.0 4- 2.2

Aurigae

TABr~S 1.

445 6296 5931 --18.9 --55.6 19.1 39.0 18.4 37.7

7450 0.34 26

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± 16.5

V V Cephei

15.00 4- 0.58 17.0: 20.0 X 108 22.6 × l 0 s

3346 -}-278 5543 5638 --1.37± 0.39

9890 0 . 2 0 0 t 0.034 346.4 q- 11.0

e Aurigae

O~

K. O. WmOHT

151

The extent of the atmospheres of these stars may vary because of the presence of largescale clouds or prominences or, possibly, because of irregular pulsations. For ~ Aurigae, Beer and Ovenden (1951) found that egress occurred 0d4 later than predicted in 1950, Larsson-Leander (1960) observed that the 1955-56 eclipse began 0.26 days earlier than predicted and the 1958 egress occurred 0.07 days later. Bappu, Dass and Viswanadham (1965) noted that during the 1963-64 eclipse the K star was 0m.08 fainter than in 1955-56 and attributed the difference to a decrease in radius of the star. Tanabe and Nakamura (1957) suggested that the duration of totality may have increased steadily from 36.d7 in 1934 to 37d5 in 1955-56. Larsson-Leander (1960) gives the length of totality as 36d9 and that of the partial phase, 1.d5; Hardorp, Herczeg and Scholz (1966) find 36.d7 and ld.5 respectively from their data. In view of the uncertainties of these quantities, we shall adopt the duration of totality as 36d8 and of the partial phase as 1.d5 in our calculations. From the photometric colours of the stars obtained from observations made during totality and outside of eclipse, approximate spectral types and absolute magnitudes of the components can be obtained. I t is found that these values agree moderately well with those found from the observed spectra. Mean values for the visual magnitudes of the components have been taken as VK = 3.~89, VB = 6.~10, with Am = 2.m21 ; in addition to observations mentioned above, results by Grant and Abt (1959), Popper (1961), O'Connell (1964b) and by van Genderen (1964) have been considered in the mean. The spectral type of the primary component is about K 3 or K 4 Ib, although Grant and Abt suggested K 4 ii; that of the secondary component is about B6 or B7 IV or v from the colours (Grant and Abt, 1959; O'Connell, 1964b; Bappu, Dass and Viswanadham, 1965) and from the spectrum (Lee and Wright, 1960; Faraggiana, 1965). From the revised equivalent widths for the ultraviolet hydrogen lines derived from Victoria plates and Petrie's (1965) revised luminosity calibration, the absolute magnitude of the B star is found to be --2~. 3. This value may be too bright since the extent of the hydrogen line wings indicates that the star is not much, if at all, above the main sequence. However, it is probably brighter than Wellmann's (1951) value, --0.mS; perhaps the luminosity --lm. 5, which is approximately the value indicated by Keenan (1963) and by Weaver and Ebert (1964) for a B6 IV star, is the best that can be suggested now. In that case according to the magnitude difference, the absolute magnitude of the K-type component is about --3m. 7 which corresponds to a type K 3 or K 4 Ib-H. The radii of the stars can be computed from the dimensions of the system and the duration of the total and partial phases--though the partial phase ~include opacity as well as occultation effects as well as the occultation, especially in the violet region of the spectrum. They can also be computed from the bolometric magnitude and effective temperature assuming that these stars radiate as black bodies. The agreement of the two methods may be considered a measure of the reliability of the dimensions, spectral types and bolometric corrections and/or the black-body approximation and model-atmosphere calculations and the similarity of the limb-darkening relative to that of the Sun. As a first approximation in calculating the diameters of the stars, an inclination of 90 ° is assumed. The notation is that used in the determination of spectroscopic binary orbits: r is the distance from the centre of mass, a is the semi-major axis, e, the eccentricity, P, the period, and K, the semi-amplitude; u = v + co, is the argument of the latitude, where v is the true anomaly and o) the longitude of periastron. The distances of the stars from the centre of gravity of the system are computed from the formula for elliptic motion, r -

a(1 -- e ~) 1 + e cosy

(1)

152

The Zeta Aurigae Stars

with u : 270 ° for the p r i m a r y star; among the systems being considered here, only in e Aurigae is the p r i m a r y star eclipsed and for it u should be set equal to 90 °. The orbital velocity of the star at eclipse is computed from 2 4Jr2aa ( 2 1) Vorb -- p2 --

,

(2)

using the r computed above. At this time the radial velocity is Vra d :

Ke cos ~o

(3)

relative to the centre of mass. Then the transverse velocity, Vtran~, which is the velocity required to derive the dimensions of the stars, is given b y Vt2rans = Vor 2 b -- V~a 2 d.

(4)

Using values for the semi-major axis of the orbits of the primary and secondary stars, it is found t h a t the transverse velocity for the K star is 28.3 km/sec at the time of eclipse, and for the B star it is 28.4 km/sec, giving a total transverse velocity of 56.7 km/sec. I t takes 38.3 days for the B star to cross the diameter of the K star, and 1.5 days for the B star to cross the limb of the K star; this latter value is relatively less certain because the limb of the K star is not well defined. As the inclination is very probably within a few degrees of 90 °, the diameters m a y be taken as 182 × 108 k m and 7.12 × l0 s k m or, in terms of the solar diameter, 130 (~ and 5.1 Q . Since the lengths of the total and partial phases are taken from the visual V data, the errors in these values should not be large since the opacity and line-absorption effects t h a t are important for the B and U observations are minimized at the longer V wavelengths. Theoretical calculations of diameters based on effective temperatures, bolometric corrections and absolute magnitudes are necessarily uncertain for the giant stars; it is thought t h a t the uncertainties for the main-sequence B stars m a y be somewhat less. The calculations are based on the formula log R ---- 8.46 -- 2 log Te -- 0"2Mbol

(5)

(Harris, 1963). Effective temperatures for the B stars are taken f r o m Harris' 0963) Table 7, and for the K stars from Keenan's (1963) Table 5. The adopted spectral types and absolute magnitudes have been discussed above, and the bolometric corrections are taken from Harris' 0963) Tables 1 and 7. Using the above data, the diameters, in terms of the Sun, for the K and B stars are 200 G and 4.1 Q respectively. If the diameters computed from the eclipse data are inserted in E q . (5), absolute magnitudes of --2~.7 and --2.~0 are obtained. These results and those for the other stars are summarized in Table 2. (fi) 32 Cygni An orbit for this star was computed by Cannon {1918) from 117 Ottawa observations; it was re-computed b y Wright (1952} who also obtained an orbit based on 96 Victoria two-prism observations. These orbits were derived assuming a period of 1140.8 days, which was based on an approximate time for mid-eclipse in 1949, and on H a r v a r d objectiveprism observations made in 1890 and 1896 t h a t showed a K - t y p e spectrum rather t h a n the usual composite spectrum. Wellmann (1957} discussed the 1952 eclipse observations and found t h a t the period between the 1949 and 1952 eclipses was 1149 days. Photo-

K. O. WRm~r

153

electric observations of the system at the phases of eclipses since that time, coordinated by I.A.U. Commission 42, have provided much additional data without, however, yielding a final answer to this and other problems. Herezeg and Schmidt (1963) have summarized the data from the 1949 to 1962 eclipses and find a period about 1147d.8, with variations of about 0d.5. Doherty (1967) has suggested that the apparent change in period between the early and the recent observations may be real, and that it may be related to a mass loss from the primary to the secondary star. However, it is very difficult to be certain that the B-type spectrum is missing in underexposed spectrograms and the early observations should be examined critically again before the large period change is accepted. Although the period change noted above does not introduce any serious error in the computed orbit, revised elements for the system were derived from the Victoria observa. tions using the period 1147.8 days, with the results given in Table 1. The changes in e and V0 are slightly larger than the probable errors but may be explained in part by the method of computation since individual observations were used in the computer-derived 1966 solution whereas normal places were employed in the 1952 analysis. From these elements it is found that the semi-major axis of the orbit of the primary star is about 2.5 × 108 km, which represents an orbit somewhat smaller than that of ~ Aurigae. A value of K s ~ 47 km/sec for the secondary star was given by Wright (1952). This value is based on only two observations and was derived before the technique of separating the secondary spectrum from the observations was refined by Wright and Lee (1959). Therefore although the agreement between the two observations is fairly good, the ratio K B / K K . ~ 2.8 is considerably larger than has been found for the other similar systems and a value of 2.0 has been adopted as probably an upper limit; thus KB ~ 34 km/sec has been listed in Table 1. The secondary spectrum of 32 Cygni should be studied further. I t has been known, almost since McLaughlin (1949a,b) announced that the system was eclipsing and that it showed atmospheric effects similar to those of ~ Aurigae, that the eclipse is nearly grazing. Assuming that total eclipse lasted 13 days, Wright (1952) found that the angle of inclination is 82°; Wellmann (1957) decided that the K star must be 350 rather than 195 times larger than the Sun and therefore adopted an inclination of 72 ° which was accepted by Wilson (1960). From data obtained at the 1952-53 and the 1955-56 eclipses Botsnla (1957, 1962) derived inclinations of 79 °.7 and 81 °.5. The extent of the atmosphere seems to vary from eclipse to eclipse as was found for ~ Aurigae. From spectrographic evidence at the eclipses of 1949 and 1952, Wellmann (1957) concluded that totality lasted at least 9 days in 1949 and nearly 12 days in 1952; similarly Scholz (1965) decided that the eclipse was total for 6 days in 1962. On the other hand, Herczeg and Schmidt (1963) considered that 1962 photoelectric measurements made with their V filter and combined with those of Chandra and Pande (1960) in 1959 indicate a total eclipse of only 1 to 3 days with their V filters, though the eclipse is somewhat longer in the B and U colours. The difference observed with the different filters is undoubtedly related to the greater atmospheric extinction in the ultraviolet produced by the additional line absorption, first noted by Roach and Wood (1952) for ~ Aurigae. The spectrographic observations, especially those made near the Call K fine, refer to a similar region of the spectrum; one criterion of totality is the reappearance of the emission at the centre of the K line, and another is the radical change in the nearby spectrum from the relatively few strong low-excitation lines, produced chiefly by FeI and Till atoms in the chromosphere, to the multitude of lines of the pure K-type spectrum. Many more observations will be required to determine the average length of totality in different regions of the spectrum. Present data suggest a duration of about 10 days for the region just below 4000 A for most eclipses. 11

B~VIA Vol. 12

154

The Zeta Aurigae Stars

The spectral type of the primary spectrum, as observed during totality, has been given as K 5 1 b y McKellar and Petrie (1958), based onWellmann's (1957) estimate. I t is certainly not much different from the other K - t y p e stars of this group. The CaI line at ~4227 and the strong Fer lines seem slightly broader and stronger t h a n the same lines in the spectra of ~ Aurigae and 31 Cygni; thus the spectral type m a y be as late as K5. The ratio SrII, 24078/Fei, ~4072 seems to be a little greater, corresponding to a luminosity class closer to I a b if the other stars are Ib. However, the differences are not great and it seems doubtful whether the absolute magnitude is more than 0m-5 brighter t h a n the other stars. Herczeg and Schmidt (1963) find t h a t their colours fit a star of type K 4 I a b and since this classification agrees well with the spectroscopic evidence, it is accepted here. For the secondary star, Wright (1952) proposed a type about B8 in his original paper where a first a t t e m p t was made to separate the spectra, but the spectrum is almost certainly earlier, and has been found to be B5 IV-V by Faraggiana et al. (1965), using a similar method, and B4 v b y Herczeg and Schmidt (1963) from a study of the colours of the system when observed in, and out of, eclipse. Chandra and Pande (1960) found a difference of 01.~025 for 32 Cygni between V magnitudes observed during totality and outside of eclipse; from these data they give m g 4.m14 and m B ~ 8m24. This large magnitude difference seems unduly large from examination of the spectra, and a difference in visual magnitude of 2:15 has been adopted; even this difference seems large since the wavelengths where the two spectra are equal (about 4000 A) are very similar in 32 Cygni and in 31 Cygni. Absolute magnitudes of My = --5.0 for the K star and --2.5 for the B star have been adopted in computing the diameters from Eq. {5). The diameters are found to be 395 Q for the K star and 6.2 Q for the B star. Since, in other systems of this type, diameters computed from the effective temperature and luminosity do not agree with those derived from the eclipse data, it does not seem justifiable to derive the orbital inclination for 32 Cygni. ~

(iii) 31 C y g n i

31 Cygni was announced as an eclipsing system by McLaughlin (1950) and he predicted an atmospheric eclipse in 1951. The star was observed intensively at Victoria and the observations have been described b y McKellar and Petrie (1958); it has now been studied over more t h a n a complete cycle there and the elements given in Table 1 are based on these observations combined with a few obtained at times of m a x i m u m velocity separation by Popper at Mount Wilson observatory (Wright and ttuffman, 1968). An earlier series of radial velocities derived from Lick observatory plates was discussed by Vinter-Hansen (1944). She used a period of 3806 days, whereas the best value, based on photometric and spectroscopic observations obtained at the 1951 and 1961 eclipses, seems to be 3784,d3 (Lindblad and Pipping, 1963; Herczeg and Schmidt, 1962; O'Connell, 1964a). Vinter Hansen's data have been re-computed using the revised period, but the elements are changed very little. However, her e = 0.131 and V0 = - - 6 . 8 7 km/sec are appreciably different from the Victoria values; the latter difference, Lick--Victoria ~ ~-0.86 km/sec is comparable to the value ~-0.70 km/sec t h a t has usually been applied to similar observations. The difference in eccentricity is larger than expected and has not been explained. When all available observations were plotted together, the principal difference between the Lick and Victoria data seemed to be the large scatter at times of m a x i m u m separation; this was ascribed at first to atmospheric motions producing rather large random fluctuations in the radial velocities, but the good fit of the individual Victoria observations

K. O. WinGer

155

relative to the mean curve, shown in Fig. 2, does not seem to require such an explanation. More observations are needed to determine whether the large-scale random motions which seem to occur in the atmospheres of W Cephei and e Aurigae can be observed also in the atmosphere of 31 Cygni. J~t 2,430,000 + 3764

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FIG. 2. Radisl velocity curves for 81 Cygni. Radial velocities of the secondary component of the system have been measured for eighteen of the best Victoria plates b y Wright and Huffman (1968), using the subtraction techniques first described b y Wright and Lee (1959). The results are not completely satisfactory because there is a difference of --4.5 kin/see between the values of the systemic velocity, V0, derived from the primary and the secondary data when the same values for e, To, and ~o (or 180 ° + ~) are used. The ratio, K~:/KB gives the mass ratio, ~/<]+ggB = 1.49, which is probably the best value t h a t can be obtained from the present data. Thus masses, ~ K = 9.2 (D and d e s = 6.2 (D are derived, and semi-major axes, a/< = 6.74 × 108 k m and aB = 10.05 × 108 km. F r o m observations obtained at the 1951 and 1961 eclipses of 31 Cygni, totality is found to last approximately 60.5 days (O'Connell, 1964a, Herezeg and Schmidt, 1962) and the partial phase about 1.9 days. From the dimensions of the orbit, these data correspond to minimum diameters, relative to the Sun, of 135 for the K star and 4.0 for the B star. These dimensions are uncertain to the extent t h a t the extinction of the atmosphere is greater in the ultraviolet region of the spectrum t h a n in the visual, and the inclination of the orbit, though nearly 90 ° , is not known exactly. These values should, however, be considered quite reliable. The dimensions can also be derived from the spectral types, and thus the effective temperatures and absolute magnitudes. As noted above, the spectrum of the primary star for 31 Cygni is almost identical with t h a t of ~ Aurigae, and was considered to be K 3 or K 4 I b b y MeKellar and Petrie (1958). Faraggiana and H a c k (1963} subtracted the 11.

156

The Zeta Aurigae Stars

K - t y p e spectrum from the composite spectrum and found t h a t the reconstructed earlytype spectrum was about B 5 v. The recent discussion of the Victoria observations by Wright and H u f f m a n (1968) classifies the spectrum as B4 and the absolute magnitude --2.m7 from Petrie's (1965) H 7 luminosity calibration. Herezeg and Schmidt (1963) derive spectral types of K 3 . 5 I b and B4 v from an analysis of the colours measured during and outside of eclipse. They (1962) find a magnitude difference between the two stars of 2 .m55. This difference, from similar measurements, was found to be 2.~47 b y Lindblad and Pipping (1963), 2m58 b y Kwee and van Genderen (1962) and 2m74 b y O'Connell (1964a); the mean difference in visual magnitude is adopted as 2m-59. I n order to calculate the size of the stars, the spectral types are assumed to be K 4 I a b and B4 r v - v , with absolute magnitudes --4m6 and --2'.n0, and bolometric magnitudes --5.~7 and --3m6, and effective temperatures of 3550°K and 17,600°K. Using these data, the diameters, relative to the Sun, are 320 Q and 4.9 Q. These results are summarized in Table 2. (iv) VV Cephei McLaughlin (1936) suggested t h a t VV Cephei is an eclipsing system somewhat similar to t h a t of ~ Aurigae although the late-type primary star is of type M2 I a according to Keenan and Wright (1957), and the secondary is probably a Be star. From H a r v a r d patrol plates, Gaposchkin (1937) derived a period of 7430 days and the numerous observations of the 1955-56 eclipse b y Larsson-Leander (1957b, 1959) and Fredrick (1960) confirm this period although there seem to be real differences between ingress and egress phases; the fact t h a t the p r i m a r y star is itself a variable m a y explain some of these differences, and irregularities in the very extensive atmosphere m a y explain most of the remainder. Peery (1966) has computed a new orbit based chiefly on radial-velocity measures of Michigan spectrograms combined with measures of Victoria plates from 1956 to 1960. He used a period of 7450 days which had been determined b y McLaughlin. Table 1 includes the orbital elements derived b y Peery for this system. There seems little doubt t h a t the masses and dimensions of the system are large, but much of the information remains uncertain. Fredrick (1960) has suggested t h a t astrometric observations show t h a t the parallax is considerably larger t h a n usually accepted, and therefore t h a t the primary star m a y not be a supergiant; the high-luminosity characteristics of the spectrum are explained b y Fred_rick as the result of the large magnetic field t h a t was observed b y Babcock (1958). Peery's radial-velocity curve for the M star shows a large scatter which is probably real and caused by large random motions in the supergiant atmosphere. The :radial velocities of the secondary spectrum have been studied only b y Peery, who based his measures on the relative positions of the emission Hfl line over the cycle. These velocities give a systemic velocity, VB, for the B star t h a t is 37 km/sec, more negative t h a n t h a t for the M star, but the variation in velocity seems to be real. The mass ratio is the :ratio of the semi-amplitudes of the range in radial velocity, KM and KB and, allowing for the V0 shift, is 2.04 according to Peery. These values, combined with the other orbital .elements give masses of 84 Q and 41 Q, assuming an inclination of 90 °. These masses a r e unusually large, and would be diminished if the mass ratio were lower, but there is no evidence for this as yet. There are, however, numerous Victoria observations of H a .obtained with high dispersion. I t is hoped to measure tracings of these H a spectrograms, a n d t h a t the velocities so obtained when compared with those of the M-type star, will give somewhat more accurate results than the Hfl data. ~ * Recently Wright and Larson (1969) have obtained radial velocities for the Ha emission line of the B star, which indicate that K B is of the same order as Kz~. Therefore it seems probable that the masses ,of the stars in the VV Cephei system are 20 to 30 Q rather than 84 O and 41 (~ as given by Peery (1966).

Spectrum Luminosity, M, Spectroscopic Eclipse data Observed magnitude difference, Am Temperature Tet~ Bolometric correction Diameter Spectroscopic Eclipse data Mass, . ~ sin i

®

®

oK

200 130 8.0 4- 1.2

4.1 5"1 5.8 4- 1.1

3550 15,400 --1.1 --1.1

--1.5 --2.0 2'21

B6 + V

K4 Ib --3.7 --2-7

Secondary

--2.5

--5-0

19:

(3O)

395

3500 --1.1

B 4iv-v

K5iab

10:

6.2

17,600 --1.6

(2.5)

Secondary

Primary

32 Cygni

--2.0 --1.6 2"59

--4.6 --2.7

320 135 9-2 ~ 2-2

4.9 4.0 6.2 + 2.5

17,600 --1.6

B4V

K4Ib

3550 --1.1

Secondary

Primary

31 Cygni

i

1,000 2,000 84

3,000 --1-2:

--6-3 --7-8

M2I~Iab

Primary

(102) 41

Be:

Secondary

VV Cephei

Luminosities, diameters and masses of the Zeta Aurigae stars

Primary

Aurigae

TABLE 2.

175 295 15.5

5 (1000) 13.7

30,000 --3:

--3:

--7.5 --8"6 7200 0

B:

Secondary F 0 lap

Primary

e Aurigae

.o

158

The Zeta Aurigae Stars

The light curve of VV Cephei has been observed primarily b y McLaughlin (unpublished), Larsson-Leander (1957a, 1959) and by Fredrick (1960), who has discussed the eclipse data, as well as McLaughlin's series of observations made between 1932 and 1957. The latter show two types of variation, one a long period of about 13.7 years and amplitude 0.~15 and the other a shorter period with average time between peaks about 349 days, and amplitude 0m.3. Fredrick discusses the different observed times of contact in blue and yellow light at the time of the 1955-57 eclipse and interprets the data in terms of the eclipse of the B star and its shell. His times of contact are: B star Shell First contact J.D. 2,435,656 J.D. 2,435,626 Second contact 5683 5713 Mid-eclipse 5931 5931 Third contact 6179 6149 Fourth contact 6206 6236 These values indicate that the radius of the shell is three times as large as that of the B star; from the Victoria spectrograms, McKellar, Wright and Francis (1957) estimated that the shell might be twice as large as the star. Peery uses nearly the same data and concludes from his orbital elements that the minimum diameters of the M- and B-type stars are 1620 Q and 88 G assuming an orbital inclination of 90 °. If the effective temperature of the M star is 3000°K, My ---- --6m29 which, in view of the uncertainties, is in fair accord with Keenan and Morgan's (1951) estimate of --Tin.0. The spectrum of VV Cephei has not yet been explained satisfactorily. In totality it has been classified as M2 + Ia-Iab. The B star shows strong hydrogen absorption lines with emission on the violet edge for the higher members of the Balmer series, though it can be seen with varying intensity on both sides of the absorption for H a and H/~. No helium lines have been detected although several out-of-eclipse plates have been compared with the reduced spectra taken during totality, using the subtraction technique of Wright and Lee (1959). Pcery (1966) has explained the reversed P Cygni profiles of the hydrogen lines as infalling streams of matter being transferred from the M star to the companion. Numerous emission lines occur in the violet and ultraviolet regions of the spectrum: [FeII], F e ~ , Tiff, Crrf, Mn~, Nil~ and Scn, as noted by Struve (1944). A line at ~3806.3 occurs in emission and Bidelman, in private correspondence, concurs that it is probably produced by [Cu~]. These lines do not v a r y with phase and probably originate in a circumstellar envelope surrounding the system. Since this system is so complicated, with a variable M star and an envelope surrounding the " B " star, which itself may vary, no attempt has yet been made to make further comparisons between the dimensions of the stars as derived from the colours, spectral types and visual magnitudes, and from temperature, luminosity and radius relations. Although a few infrared spectrograms have been obtained (Hynek and Keenan, 1945; Keenan and Hynek, 1947; Keenan, 1964) which show strong 0 1 lines at certain phases and are almost certainly related to the secondary eclipse*, details concerning the spectrum and energy distribution of the early type star are almost unknown. The variation with wavelength of c~ = ICB/IcM, the ratio of the apparent continua of the two stars has been obtained by Peery (1966); it varies from 0.3 at 4300 A to 1.0 at 4092 A. A preliminary value obtained from the Victoria high-dispersion plates ranges from 0.3 at 4400 A to 1-0 at 4000 A and 4.0 at 3800 A; the results depend on how the continuum is drawn and, * Glebocki and Keenan (1967) observed these lines again in 1964 and 1965, thus conforming these effects at secondary eclipse.

K. O. W~OHT

159

for VV Cephei, great care must be taken that no emission lines are confused with peaks in the spectra where few, if any, absorption lines occur. The large increase in c~ at the shortest wavelengths has been observed also for the other ~ Aurigae stars and is related both to the increase in intensity in the continuum of a hot star and the decrease for a cool star, as well as to the overlapping of the numerous atomic and molecular absorption lines, and possible additional sources of continuous absorption in the cool star. (v) ~ Aurigae This system, which has the longest known period of any eclipsing spectroscopic binary, 27.08 years, has a fiat minimum with a depth of 0.ms which lasts 330 days, and a partial phase of 357 days. I t has been an enigma since the eclipse of 1927-29 because the spectrum during minimum is very nearly the same as that of the principal component, F 0 Ia, and remains single throughout the cycle whereas the depth of minimum implies that the components are almost equally bright. Kuiper, Struve and StrSmgren (1937) explained this phenomenon b y considering that the companion is semi-transparent to the light of the primary star during eclipse because the orbital inclination of the system is about 70 ° and therefore its radiation passes through the outer layers of the companion. Struve (1956) modified this hypothesis by suggesting that both stars are surrounded by nebulosity filling the Lagrangian loops ; the companion is unknown since only its outer shell produces the eclipses. The additional lines produced before and after totality are originated in gas streams flowing from the shell towards the principal star. Miss Hack (1961) studied spectra of ~ Aurigae obtained at Mount Wilson during and after totality, and concluded that the companion could be a hot B-type star with an extensive shell which produces the opacity necessary for the eclipse. Observations of H a near eclipse tend to confirm this hypothesis which was, indeed, suggested by Wright and Kushwaha (1957) when they discussed the Hc¢ observations. These proposals have been considered and expanded by Morris (1963, 1965) who thinks that the system originated as a pair of O-type stars with a mass-ratio somewhat greater than the present value of 1.13. The more massive star has evolved into an F supergiant and is undergoing mass ejection through the Lagrangian point, while the less massive star has remained close to the main sequence. Thus part of the ejected mass remains surrounding the F-type star; the remainder passes through the inner Lagrangian point and forms a disc around the less massive star. The inner region of the disc is totally ionized by the central star; the intermediate region is heated sufficiently to become nearly opaque near the orbital plane, and the outermost zone remains transparent and produces the shell spectrum observed during eclipse. The orbital elements of this system were first determined by Ludendorff (1924) based chiefly on Potsdam spectrograms. There is a large scatter among the radial velocities measured in a given year, but the averages fit the curve fairly well; Kuiper, Struve and Str6mgren (1937) calculated a new orbit which changed T from 1920.6 to 1924.2, but the other elements were similar. Morris (1962) combined these data with more recent Mount Wilson, ¥erkes and Victoria observations. The elements given in Table 1 are effectively the same as the 1962 solution but have been derived from a slightly revised computer programme after five iterations. The principal difference between these elements and those obtained by Kniper, Struve and Str6mgren (1937) is that the eccentricity has been decreased by a factor of two. There remains a very considerable scatter of the radial velocities about the mean curve; it seems probable that the scatter is real and that the velocity spread is the result of random motions in the stellar atmosphere. The semi-major axis is 1.97 × 109 km if an orbital inclination of 90 ° is assumed. Morris (1965) derived

160

The Zeta Aurigae Stars

an absolute magnitude for the F star of --71~25 from the strength of the interstellar lines and from the astrometric orbit. Theoretical models suggest a mass of 15.5 Q and a radius 174 Q. From the mass function, the mass of the secondary star is 13.7 Q, and from the eclipse data its radius is 850 Q. Numerous photoelectric observations of e Aurigae were obtained at the time of the 1955-57 eclipse. The most detailed discussion has been given by Larsson-Leander (1959) who combined his data with those of Gyldenkerne to derive the following times of contact : first, J.D. 2,435,312 ; second, J.D. 2,435,442 ; third, J.D. 2,435,800; and fourth J.D. 2,435,965. The duration of eclipse, 653 days, is less than the mean of previous eclipses, 714 days, and the duration of each phase seems to vary from eclipse to eclipse. However, the period, between 9889 and 9906 days, depending on the time of contact, and the phase considered, agrees well with a study of earlier eclipses (9890 days) by Gfissow (1936). No further discussion of the light curves is given here since they are studied in detail by Gyldenkerne (1970).

THE OUTER ATMOSPHERES

Since the time of the 1932 eclipse of ~ Aurigae it has been known t h a t the K line of ionized calcium, 43933.66, in the spectra of this group of stars becomes enhanced well before and after the eclipses; the H line behaves in a similar fashion but the He line, only two angstroms away, makes the interpretation of the observations more difficult and therefore the K line is usually the only line studied in detail. These observations have been interpreted as absorption of the radiation of the B star b y the outer atmosphere of the primary star. Early observations seemed to inticate a regular increase in the absorp. tion intensity of the K line with time, as totality was approached; this would correspond to a rapid, but regular, increase in the density of the atmosphere. However when the intensity of the K line in the spectrum of ~ Aurigae was plotted against the distance of the B star from the limb of the K star, large differences were found beetwen different eclipses and even beetwen ingress and egress at the same eclipse. The same phenomenon has now been found in other systems. Spectra taken at the 1951 eclipse of 31 Cygni showed that the K line was not always a single line, but, at certain phases consisted of sereval, often well-separated, components. Much of the observational effort during recent eclipses, particularly at Victoria, has been directed to the detection of these satellite lines in order to determine the duration of these lines which, presumably, are produced by clouds, prominences or streams of gas in the outer atmosphere moving with different velocities. The observations of this phenomenon will be discussed for each star in turn.

(i) ~ Aurigae As noted above, the chromospheric K line (as we shall call it) was observed in 1932 by Guthnick and Schneller (1935), in 1934 by Beer (1934) and by Christie and Wilson (1935), in 1937 by Beer (1940), in 1939-40 by Wilson (1948), in 1947-48 by Welsh (1950), in 1953 and 1955-56 by McKellar and Butkov (1957) and in 1963-64 by Odgers and Wright (1965), by Hardorp, Herczeg and Scholz (1966) and by Faraggiana (1965). Welsh discussed the extent of the atmosphere at eclipses up to 1947 and Odgers and Wright considered the data for the later eclipses. The observations of the intensities of the K line at ingress

K. O. WRm~T

161

and egress for the different eclipses are shown in Fig. 3, and the general trend of the data is noted in Table 3. I t is seen that the extent of the atmosphere has varied from the very great extension observed by Beer during the 1937 ingress to the very similar, apparently normal ranges in 1934, 1939-40, ingress in 1953 and egresses in 1947-48 and 1963-64. From Table 3 it would seem that egress is nearly always " n o r m a l " and that the principal differences occur during ingress. Although Guthnick and Schneller did not make measurements on their spectra, the K line could be observed about 45 days before mideclipse which, for the dispersion used, would seem to be comparable to the 1947-48 ingress as observed at Victoria• Welsh suggested that it might be possible to relate the abnormal extensions to some permanent feature, and therefore to the period of rotation of the star (tentatively given as 785 days by Christie and Wilson, 1935). This hypothesis seemed somewhat plausible in view of the 1932 ingress data, but the numerous normal eclipses make this period unlikely. Since periastron occurs about 70 days after mid-eclipse, if it were a gravitational effect of the secondary star on the atmosphere of the primary, it would be more reasonable to observe the extension at the egress phase. I

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162

The Zeta Aurigae Stars TABLE 3. The varying extent of the atmosphere of Zeta Aurigae Eclipse

Ingress

1932 1934 1937 1939-40 1947-48 1953 1955-56 1963-64

Normal Very extensive Small Extensive Normal Normal Normal

m

Egress

Extensive Normal Normal Normal Normal Extensive Normal

Although the K-line intensity observations have been made at several observatories and the method of reduction may be somewhat different--~for example, the Victoria observers usually make some allowance for the double emission peak near the centre of the line when the line is relatively weak, several weeks from totality it seems quite definite that the extent of the atmosphere, as measured by the K-line intensity, does vary from eclipse to eclipse, and that it is different at ingress and at egress. The observations available have been plotted in Fig. 3, and mean curves for the different eclipses have been drawn; smooth curves have been adopted except where high-dispersion observations indicate large changes in relatively short intervals. At ingress, the extent of the atmosphere was probably less than normal in 1939-40 and greater than normal in 1947-48, with a much greater extension in 1937 ; the data for the remaining eclipses that have been observed could probably be called " n o r m a l " atmospheres. At egress, the atmospheric extent of the 1955-56 eclipse seems to have been greater than normal, and the remaining data can, perhaps, be combined, at least for the first 10 days after the end of totality. The difference between the 1947-48 data and those of 1963-64 might be related, in part, to the higher dispersion of the later Victoria data; however, the Merate spectra, which show the same general trends, arc somewhat comparable to the 1947-48 Victoria data and therefore the differences are considered to be real. There have been relatively few observations of satellites or component lines of the K line in the spectrum of (Aurigae. Those that were detected by Odgers and Wright (1965) on the 6 A/mm plates taken with the Victoria 48-inch telescope at the coud~ focus were quite weak and the distribution of observations was such that they usually appeared as asymmetries at the edge of a fairly strong principal component. Little information is yet available concerning the lifetimes of these satellites, and the observations of 1963-64 were too scattered to draw any conclusions. There is little doubt that there are condensations in the atmosphere of $ Aurigae and that they extend several radii from the edge of the star as defined by the eclipse observations. As will be noted for 31 Cygni, the K lines seem to vary throughout the cycle. The chromospheric lines may becomes moderately strong a few months before eclipse and then decrease in intensity before becoming strong again as the light from the B star passes through the denser atmosphere of the K star before the final extinction effects become important (Roach and Wood, 1952), and total eclipse begins. (ii) 32 Cygni Spectrographic eclipse observations of the K line have been made of 32 Cygni since 1949, especially at Victoria (Wright, 1952; Wright and McDonald, 1959; Wright and Hesse,

K. O. WmaHT

163

1969), at Hamburg, Michigan and Toronto (Wellmann, 1957; Scholz, 1965) and Merate (Faraggiana et al., 1965). The observations outside of totality are shown in Fig. 4. Most of these were obtained with dispersions of 20 A/mm or less, and therefore there was little probability that weak satellite lines could be observed. For observations obtained more than about 40 days from mid-eclipse the double emission, which is a well-known phenomenon of the out-of-eclipse spectrum, appears and is not taken account of in the same way by all observers. For most of the Victoria observations a mean out-of-eclipse profile of the emission is adopted and fitted to the intensity-continuum that has been adopted as the sum of the continuum of the B star and the absorption Call K line of the K-type spectrum (Wright, 1952). The correction for this emission usually adds about 150 mA to the intensity of the chromospheric K line, and therefore is not important within about 40 days of totality; the effect is usually omitted when the llne becomes wide and the emission cannot be detected. Wellmann (1957) has discussed the variability of these emission profiles, especially out of eclipse. They do vary but, for the eclipses of 1949 and 1952, they seem to be similar at the same phase. I t seems almost certain that the principal double-emission feature appears in the K-type spectrum; it is similar to that observed in the spectra of other K-type stars and has been used by Wilson and Bappu (1957) as a criterion of luminosity in late-type stars. Wellmann considers that the intensity changes are the result of absorption outside the atmosphere, which, because of the high inclination of the orbit, are not projected onto the B star at any time. Further work would seem to be desirable. The duration of totality and its possible variation have been mentioned in an earlier section. Wellmann found that totality lasted about 12 days for the eclipses of 1949 and 1952, but Herczeg and Schmidt (1963) in their discussion of the photoelectric observations questioned whether it lasted more than 3 days in 1959 and 1962. Figure 4 seems to show t h a t the chromospheric K-line absorption at any given phase, and especially as the projected distance of the B star from the limb of the K star became small, was less at the 1962 and 1965 eclipses than the absorptions observed earlier. The 1962 data seem to confirm the Merate observations (Faraggiana et al., 1965). The full effect of the emission has not yet been considered for the 1965 observations. Spectrographic observations obtained during totality are insufficient to determine the duration of these eclipses. The definition of totality that has been usually accepted is the reappearance of the double emission at the centre of the very strong K absorption line, as noted for ~ Aurigae. Victoria and Merate observations confirm that the eclipse was total from June 3 to 6, 1962; Herczeg and Schmidt's photoelectric observations place mid-eclipse as May 31.5 (J.D. 2,437,816.0), but no spectrographic observations seem to have been obtained before mid-totality. Assuming a period of 1147.8 days, mid-eclipse in 1965 occurred on J.D. 2,438,963.8 (July 22). Victoria spectra of July 15 and 18 indicate that the eclipse was then total; a spectrogram taken on July 27 at a dispersion of 10 A/mm also shows the characteristic emission feature at the bottom of the K line; therefore it is considered that the system was then in total eclipse. A length of totality of 12 days, or more, similar to 1949 and 1952, is suggested*. Thus the difference between the weaker chromospheric lines observed in 1962 and 1965 and earlier is not solved. I t could, however, be related to the fact that the eclipse is nearly grazing and that the limb is not sharp; hence the radiation absorbed in the deep partial phases of the eclipse, very close the limb, is very little different from the normal K-type spectrum. Future eclipses should be observed with both spectrographs and photometers in order to determine the length of totality precisely. * High-dispersion spectra obtained at Victoria during the 1968 eclipse showed a K-line in emission for 19 days, thoug it is doubtful whether the st~r shauld be considered in total eclipse during all of this time (Wright and bforbey, 1969).

164

The Zeta Aurigae Stars

Equivalent-width measurements of the K line in the spectrum of 32 Cygni have been plotted against phase from mid-totality in Fig. 4; a period of 1147.8 days from the zero phase, J.D. 2,434,374.3, adopted by Wright and McDonald (1959), has been used. Curves showing the trend of the observations for the 1949, 1952, 1959, 1962 and 1965 eclipses have been drawn and, as for ~ Aurigae, smooth curves have been drawn except where definite large changes have been observed. Differences from one eclipse to another, as shown by these curves, are not as pronounced as for ~ Aurigae, but during ingress in 1965 the e x t e n t of the atmosphere seems to have been less that at previous eclipses; a detailed analysis of the 1962 and 1965 Victoria data is being prepared by Wright and Hesse (1969). During egress, the atmospheric extent was less at both 1962 and 1965 eclipses than had been observed previously. The 1962 data may be related to the short duration of totality noted by Herczeg and Schmidt (1963). The irregularities in the 1965 curves are almost certainly real in most cases, since most of the Victoria plates are well exposed and some satellite lines can be seen easily on the intensity tracings shown in Fig. 5 and in the spectra of Plate I. Very few radial-velocity measurements of the chromospheric lines have been made for 32 Cygni. Faraggiana, G6kg6z etal. (1965) measured nineteen plates taken at the w (AJ

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166

The Zeta Aurigae Stars DAYS) d-eclipse

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167

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tracings have been rectified relative to the continuum of the B star plus the b o t t o m of the absorption K line in the K - t y p e star; the intensity of the latter changes quite slowly in the region under discussion, and has been drawn on the logarithmic tracings before rectification. Possible satellite lines have been sketched in as dotted profiles. I t is not claimed t h a t all of these lines are real, b u t wherever a clump of grains t h a t might be related to an absorption line appeared on the tracing, it was drawn as a line. The broad central features were also treated as several individual lines until definite wings showed t h a t the feature should be treated as a single unit. The centres of these lines were measured on the tracings and the measures converted to wavelengths assuming t h a t the F e I lines of R.M.T. multiplet 43 were displaced from their laboratory wavelengths b y the amount required b y the orbital velocity of the star. I t is planned to measure these lines in the usual manner, relative to the standard comparison lines, but it has not yet been possible to make this correction. The wavelengths of a representative selection of these lines have been plotted against phase in Fig. 6. The sizes of the dots in the figure are an indication of the relative intensities. The principal K line is shown with the largest dots in the diagram. I t is seen t h a t the velocity (which is directly related to the wavelength) changes rather erratically with time; the thin connecting lines represent one possible interpretation of its behaviour. The trends of the other features are even more speculative, b u t are indicated as a basis for discussion.

The Zeta Aurigae Stars

168

I t is seen t h a t m a n y of the features seem to be related and to continue with the same general intensity and at approximately the same position for periods ranging from a few days to several weeks. I n some cases the variation in wavelength is somewhat similaI to t h a t of the principal K line. However, it is not essential t h a t the velocities be the same during the lifetime of the feature since velocities of the solar prominences (which are the only other known stellar features t h a t can be related to those discussed here) certainly do not remain constant as they soar above the solar photosphere. I t seems rather unlikely t h a t a single condensation could be observed over distances comparable to the diameter of the primary star unless they are very extensive, or unless they are moving in such a way as to follow the p a t h of the projected beam from the B star. No conclusions concerning these observations seem to be justified at the moment, but they indicate the need for making similar, and even more detailed studies of chromospheric spectra at future eclipses. The last eclipse of 32 Cygni occured about J.D. 2,440,111•6 (Sept. 12, 1968). One further point might be mentioned concerning the velocity of the principal K line indicated in Fig. 6. Prior to eclipse the general trend is for its velocity to be positive relative to the Fez lines; after eclipse the trend is to negative velocities. This is different from observations made at previous eclipses, and must be checked when true radial velocities are derived. -70 1

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K. O. WnIOHT

169

(iii) 31 Cygni 0 n l y two eclipses of 31 Cygni have been observed in sufficient detail to study the structure of the K line and the effects produced by the passage of the radiation of the B star through the outer atmosphere of the K star. In 1951 the best series of observations were obtained at Michigan and at Victoria, but more details could be obtained from the latter, since the dispersion was considerably higher. The principal study of the K line was made by McKellar et al. (1959), although the important discovery of multiple K lines was announced earlier (McKellar et al., 1952). Studies of the inner chromosphere have been made by Larsson-Leander (1957a), based on Michigan spectrograms, and by Wright (1959); they will not be discussed here since the results, although based on excellent spectrographic material, are not considered definitive until more theoretical studies of the structure of the chromosphere are made. They are quite similar, however, to those obtained by Wilson and Abt (1954) and discussed by Wilson (1960) for ~ Aurigae and 31 Cygni.

PLATEII. Spectra of 31 Cygni obtained at the 1951 eclipse, showing satellite lines.

The radial velocities of the chromospheric lines have been discussed, for the 1951 eclipse, by McKellar et al. (1959) and, for the 1961 eclipse, by Wright and Odgers (1963) and by Faraggiana and Hack (1963). In 1951 the trend of the velocities relative to the orbital velocity was rather erratic but more negative than positive before eclipse. This was quite marked for the K line, especially about 50 days before totality. After totality the K-line velocities were more positive than the orbital velocity for nearly a month and then the line split into two components, one of which gave velocities approximately 20 km]sec more positive than the orbital velocity and the other gave velocities about 15 kin/see, more negative. The same general trend was observed for the Michigan plates although, when the lines were measured as double, the velocities were usually somewhat more negative than the Victoria velocities. From the velocity-graphs published by McKellar et al. (1959) shown in Fig. 7, it would seem quite possible that the two principal satellite K lines may have been produced by radiation passing through two very extensive 12

B-VIA Vol. 12

The Zeta Aurigae Stars

170

clouds during the 40 days t h a t they could be seen separately. However, the intensities showed rather large changes (in Plate I I note the red component on the plate of December 7) ; the shapes of the lines also v a r y by more than a factor of 2. Hence it is questionable whether a single cloud explains these observations t h a t cover more t h a n a month. An interesting feature of the spectra shown in Plate I I is the progression of velocities of the satellite line from October 26 on successive days until November 1, 1951. I t seems doubtful whether the negatively displaced feature of October 26 arises from the same source as those of later date, but it does seem possible t h a t the regular progression of displacements from October 27 to November 1 could be the result of radiation passing through a cloud which first is moving away from the observer and then coming back towards him. The principal line shown in this series has decreased in intensity and width quite markedly

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K. O. WRIGHT

171

during the 5 days, but the satellite, which is indistinct in the reproduction though clearly visible on the original plate, remains nearly the same intensity, although it seems broader on the final spectrogram. At the 1961 eclipse of 31 Cygni, during ingress the velocity of the K line, shown in Fig. 8, follows the orbital velocity until the line of sight of the B star is about 3.5 stellar radii from the limb of the K star; it is observed at a velocity about 10 km/sec negative relative to the orbital velocity while the B star travels more than half a radius, comes back nearly to the orbital velocity, and then remains about 15 km/sec negative almost until second contact. At egress the results measured by Faraggiana and Hack (1963) follow a similar negative trend. The observations of the satellite lines observed on the Victoria spectrograms are somewhat inconclusive in determining the lifetimes of the features. A study of the radial-velocity plot seems to indicate t h a t the features may last as independent entities for 20 to 40 days, but large gaps in the observational material make it difficult to correlate the data. The intensity data for the principal component of the K line for the two eclipses t h a t have been observed are shown in Fig. 9. The general trend of the increase of intensity as eclipse is approached is quite similar for both ingress and egress at both eclipses. The data combine the observations made at Victoria and Michigan in 1951 (McKellar et al., 1959), and, in 1961-62, at Victoria (Wright and Odgers, 1963), at Merate (Faraggiana and Hack, 1963) and at Hamburg (unpublished data sent to the author by Groth, 1964). The scatter of the data is rather large and it is difficult to draw positive conclusions about the extent of the atmosphere from two eclipses. A comparison of the intensity measurements made at Victoria and at Hamburg from the 1961 data indicates that the measured Hamburg intensities are much greater than the Victoria measures until the lines become quite strong during the inner chromospheric phase (2 weeks before second contact). This effect has been allowed for in the diagram by multiplying the Hamburg intensities by a factor 0.6 for values less than 1.0 A, and 0.8 for values greater than 1.0 A. l~ocorrection has been applied to the Merate observations. The differences between ingress and egress and between the two eclipses are considerably less than the range observed for the eclipses of ~ Aurigae. I t is thought that the extent of the atmosphere of 31 Cygni was similar in 1951 and in 1961. However, there is no doubt that the atmosphere is not uniform. Although the numerous satellite lines measured by Wright and Odgers may not all be real, many of them are sharp, clear-cut features, and a number of them have been measured also by Faraggiana and Hack (1963). Representative intensity tracings of Victoria spectra in this region are shown in Fig. 1O. Perhaps the most interesting observations made at Victoria were on three successive nights 167 to 165 days before mid-eclipse (August 7 to 9), when the projected distance of the B star was more than two stellar diameters from the limb of the K star. The satellite lines had velocities of --50, --54 and --60 km/sec on these days and it seems almost certain that they were produced by the same condensation in the atmosphere. As the series of observations progressed, satellite lines of definitely measurable intensity occurred at irregular intervals. However, for a few days, when the projected distance was one stellar diameter, even the principal K line became very weak as if the atmosphere became either very tenhous or some continuous veiling obscured the normal features. The characteristic square feature of the principal line appeared about one radius from the limb, but it became apparent that this profile was produced by several components with different velocities, when, after a few days, the line became narrower again. About one-haif a radius from the limb, the line began to show fairly symmetrical wings, when the region of the inner chromosphere was reached. 12"

The Zeta Aurigae Stars

172

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174

The Zeta Aurigae Stars

McKellar et al. (1959) have discussed the profiles of the K lines observed at the 1951 eclipse and found t h a t the narrowest lines correspond to turbulent velocities of about 10 kin/see and the broadest lines correspond to velocities of more than 40 km/sec. The observed satellite lines indicate t h a t the broad lines m a y be the result of two or more clouds in the atmosphere moving with speeds differing b y this amount, but the narrow lines probably represent small random motions in the atmosphere as a whole. (iv) VV Cephei P a r t l y because of the complexity of the system and the m a n y different features repre. ted on the spectrograms, a detailed analysis of the outer atmosphere of VV Cephei has not yet been made. Studies of the inner chromosphere have been made b y Gaposchkin (1937), Goedicke (1939) and Peery (1966). While discussions of the inner chromosphere are not included in this paper, it is worth mentioning t h a t turbulent velocities of 48 km/sec were derived b y Peery for the chromospheric lines. Discussions of the Victoria observations of the 1956-57 eclipse have been given b y Wright and McKellar (1956) and b y McKellar, Wright and Francis (1957). The most interesting feature of the chromospheric observations is the appearance of double lines in the chromospheric spectrum for several weeks, ending about a month before the beginning of the partial phase*. The separation of the lines varied with the element producing them, but was about 25 km/sec for the iron group. The separation was greater for lowexcitation TiII lines, but it seems probable t h a t the latter are formed, at least in part, in the circumstellar envelope surrounding the system, which makes the interpretation of the observations even more difficult. However, it is probable t h a t two separate " c l o u d s " of material were between the B star and the observer when these observations were made. Numerous observations have been made at Victoria since 1956. A brief account of the d a t a and proposals for their analysis has been given b y Wright (1967), and little more can be added at the present time. The K line has not been studied in detail, but two separate absorption features can be seen on most well-exposed plates. Deutsch (private communication) has suggested t h a t b o t h m a y be related to m a t t e r in and around the system, although it would be useful if an interstellar line could also be detected in view of the divergent results obtained for the distance of the system (Fredrick, 1960)**. (v) e Aurigae " C h r o m o s p h e r i c " lines, in the sense they have been defined in this paper, were first observed in the spectrum of e Aurigae in 1930 b y Adams and Sanford (1930}. K r a f t (1957} has made an analysis of these observations. For the 1955-56 eclipse the best series of high-dispersion spectrographic observations were obtained b y Struve at Mount Wilson and b y Wright at Victoria. Prior to the eclipse Struve (1956} proposed a new model for the system to explain asymmetries and other features t h a t had been observed ; it involved nebulosities surrounding both stars and flows of gas between the stars. At the time of the eclipse Struve and Pillans (1957) observed double and, in a few cases, triple chromospheric lines, which must arise in different regions of the system; these observations lend some * A review paper on "The VV Cephei Stars" has recently been published (Cowley, 1969). High-dispersion observations of the spectrum of VV Cephei are being continued at Victoria. Changes in both emissio~ and absorption features, particularly in the far ultraviolet region suggest that an envelope surrounding one or both stars, and perhaps even the system, must be very extensive. ** The Ti I I lines, 3759 and 3761 A, were also observed as double in 32 Cygni at the 1968 eclipse, from 31 to 22 days before predicted mid-totality.

K . O. WRIGHT

175

support to Struve's model. Struve, Pillans and Zebergs (1958) analysed the radial velocities of ~ Aurigae, especially these measured on plates taken near eclipse, and Wright (1958) measured most of the Victoria plates and showed how the chromospheric lines were displaced nearly 20 km/sec positive before eclipse and about 30 km/sec negative after eclipse, relative to the velocities of the background lines; these results are shown in Fig.11. Miss Hack (1959) studied a few of the Mount Wilson plates just before and just after the 200 r

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end of totality in 1956; differences in the physical conditions of the atmosphere derived ~rom these observations are small. Morris (1963) also studied a few of the Victoria plates taken at the time of the 1955-56 eclipse and found a dilution factor of about 0.1; his general conclusion was that the secondary component is a hot, massive star surrounded by a disc of gas lying in the orbital plane. This disc has probably been accretod from material ejected from the primary star; the opacity of the disc which produces the eclipse can be explained as being due to the continuous absorption of hydrogen.

17 6

The Zeta Aurigae Stars DISCUSSION AND SUMMARY

I n this article, most of the principal data concerning the elements of the orbits, the dimensions and the outer atmospheres of the giant components of the Zeta Aurigae stars have been considered and summarized. The data are collected in Tables 1 and 2. Most of the following discussion relates to ~ Aurigae and 31 Cygni since the data for these stars are more reliable t h a n those for the other stars. The masses of these stars seem to be fairly well determined, but their diameters and luminosities are less well known. The masses of the K stars are about nine times t h a t of the Sun; the data for 32 Cygni are insufficient and, unless this system is quite .different from the other two, it seems probable t h a t a lower mass ratio t h a n has usually been accepted will be found when further studies of the secondary spectrum are made. The masses of the B stars are about six times t h a t of the Sun, which seem to be reasonable for the main-sequence stars t h a t these appear to be. The diameters of the K - t y p e components of ~Aurigae and 31 Cygni derived from the eclipse data agree very well even though the periods differ by a factor of three; therefore it is thought t h a t the values of 130 solar diameters are quite well determined. These diameters are based on visual data, and therefore extinction effects observed at the time of eclipse should be minimal. The data for the B stars are less certain because effects at the limb of the K star will affect the values found for the smaller B star to a greater extent; however, the values of four to five times the solar diameter seem to be satisfactory. These are minimum values since an inclination of 90 ° has been assumed, but it is very unlikely t h a t the results will be increased by even 10 percent since the eclipse of 32 Cygni seems to be very nearly grazing. For this system the effects at different eclipses v a r y quite markedly, and they are different from those observed for ~ Aurigae and 31 Cygni (see McLaughlin, 1952). Thus it would seem t h a t the eclipses of ~ Aurigae and 31 Cygui are nearly central. The luminosities of these stars are difficult to determine. I t might seem t h a t the application of Petrie's (1965) H~-luminosity calibration would give reliable luminosities, but the separation of the B-type spectrum from the many-lined K - t y p e spectrum is difficult and at H 7 the B spectrum is so weak t h a t it is impossible to obtain useful information for either radial-velocity or intensity measurements. For the ultraviolet hydrogen lines from ~3770 to ~3970 measurements can be made, but the scatter of the radial velocities of the secondary component, shown in Fig. 2, shows t h a t high accuracy cannot be exp e c t e d - t h o u g h even in unblended B-type spectra the broad lines make the measurements much less certain t h a n those obtained for a K - t y p e star. For each of these B-type stars the luminosity based on hydrogen-line equivalent widths is greater t h a n has been accepted as normal for main-sequence stars of the observed spectral type (Keenan, 1963; Blaauw, 1963; Weaver and Ebert, 1964); it is, of course, difficult to determine the luminosity class on the MK system from the short range of spectrum in the ultraviolet t h a t can be obtained using subtraction techniques. The ratio of the continua of the two stars, ~ ~ I t B / I c K , appears implicitly in the derivation of the equivalent widths of the lines in the B-type spectrum; these values would be increased, and the luminosity correspondingly decreased, if the contribution of the K star were greater t h a n the adopted value. For ~ Aurigae, Lee and Wright (1960) found t h a t the intensities of the continuum of the two spectra were equal (~ ~ l) at 4110 A, and for 31 Cygni, Wright and Lee (1959) found t h a t they were equal at 3875 A. These results were obtained from high-dispersion spectrograms and the ~'s so derived are lower, corresponding to a greater contribution from the K star t h a n similar measurements made on lower-dispersion spectra (Hardorp, Herczeg and Scholz, 1966). Thus there seems to be little justification for decreasing the luminosity of the B stars b y any large amount.

K. O. WRIGHT

177

As indicated earlier, the relation between luminosity, effective temperature and diameter of the stars depends on the assumptions relating the effective temperature to the total radiation of the star and, although the bolometric corrections have been derived from model-atmosphere calculations for main-sequence stars, these corrections are quite uncertain for giant stars. The formula assumes that conditions in the star, including limb darkening, are similar to those in the Sun but the extensive atmospheres of the giant stars must produce quite a different distribution of energy across the stellar surface. However, this relation [Eq. (5)] is the only one available for comparing stellar diameters with other parameters. As shown in Table 2, the luminosities a d o p t e d for the K-type stars give diameters about twice those computed from the eclipse data. The luminosities of the B stars were adopted after considering the photometric and spectroscopic data, and the luminosities of the K stars were derived from the observed V magnitude difference between the eclipsed star and the out-of-eclipse data. This observed magnitude difference seemed to be too large for 32 Cygni, but the values for ~ Aurigae and 31 Cygni seemed reasonable. However, when the calculations were worked through in the reverse direction, beginning with the diameters calculated from the eclipse data, the absolute magnitudes of the K stars are at least a magnitude less than those derived from the B-star luminosities, and the magnitude difference is decreased by more than a factor of 2. I t would seem that the best estimate of the absolute magnitudes of these stars is, perhaps, a mean of the spectroscopic and eclipse data. I t is of interest to consider the absolute magnitudes of the K stars in view of the recent discussions of Hodge and Wallerstein (1966) and of Wilson (1967) relative to Wilson and Bappu's (1957) relation between absolute magnitude and the widths (log We) of the K.line emission feature in late-type spectra. Hodge and:Waller~tein concluded that the luminosities for a given log W e should be increased for the brighter stars in order to fit observations of stars w i t h well-determined trigonometric parallaxes. By forcing the Hyades giants to fit their new calibration, Hodge and Wallerstein claimed t h a t cluster ages could thus be decreased enough to obtain agreement between the age of the oldest clusters and the expansion age of the Universe. On the other hand, Wilson rederived his calibration using luminosities derived from sixty-five stars with good trigonometric parallaxes, and found that the Hyades stars fit this calibration well, and that it is very nearly the same as the Wilson (1959) Sun-Hyades calibration; thus a change in the modulus of the Hyades by 0.m4 as suggested by Hodge and Wallerstein is not required. The relevance of the above discussion to the $ Aurigae stars is that, in the Wilson calibration, the fit of ~ Aurigae was considered a verification of the calibration when Wellmann's (1951) luminosity estimate of --2.m2 was used for the K star; on the SunHyades calibration, log W 0 ~ 2.0 for ~ Aurigae corresponds to a luminosity --2m. 3. If a higher luminosity (say --3-m2) is adopted for the K star of ~ Aurigae, the deviation of this star from the Wilson-Bappu calibration is increased and, indeed, a luminosity of --3.m2 fits Hodge and Wallerstein's curve quite well. For 31 Cygni, McKellar and Petrie (1958) obtained a value of log W 0 = 2.08 from Victoria spectrograms, which corresponds to My = --3.m4 on the Wilson Sun-Hyades calibration and --4.m5 on the Hodge-Wallerstein scale. McKellar and Petrie adopted a value of --4m. 0 for the K component of 31 Cygni; our value is --4.~6 from the MK calibration and --2~.7 from the eclipse data, with a mean of --3.m6. The high-luminosity portion of the Wilson Sun-Hyades curve was tested by the red supergiants in h and Z Persei and Wilson and B a p p u stressed the point that their M y - log We calibration was linear over a range of fifteen magnitudes. The present analysis does not add significantly to our knowledge of the luminosities of the K - t y p e supergiants. The mean of the luminosities listed for $ Aurigae and 31 Cygni in Table 2

178

The Zeta Aurigae Stars

places t h e m above the Wilson calibration curve and slightly below t h a t of Hodge and Wallerstein. No positive conclusions can he drawn from this and the uncertainties in the luminosities of these stars must be at least 01.~5. Additional studies of the K stars in these systems are needed. I n this paper only a few more comments will he made concerning Faraggiana and H a c k ' s (1966) comparison of the spectra of the K - t y p e components of ~ Aurigae, 31 Cygui and 32 Cygni, which was based on Merate 22 A/ram plates taken during total eclipse. They found t h a t the line intensities in the spectra of ~ Aurigae and 31 Cygni are fairly comparable, hut the intensities of the same lines in 32 Cygni are greater; curves of growth gave turbulent velocities of 7.6, 8.3 and 12.0 km/sec, respectively for the three stars. Groth (1955) derived a velocity of 14.5 km/sec for ~ Aurigae, and Wright (1957) gave a preliminary value of 7.5 kin/see for 31 Cygni. Faraggiana and H a c k also derived colour temperatures for the K - t y p e stars b y comparing the observed ~ ( = ICe/IcK ) values and adopting temperatures for the B-type components. The derived temperatures for the K stars were as much as 1000°K lower than the excitation temperatures of 3600°K t h a t were f o u n d ; t h i s excitation temperature m a y be compared with the value of 3700°K found b y Groth for ~ Aurigae, and t h a t of about 4()00°K obtained by Wright for 31 Cygni. Undoubtedly on low-dispersion plates the continuum of a K - t y p e star is drawn progressively lower from violet to ultraviolet regions of the spectrum as a result of the increasing number of overlapping absorption lines, and this effect corresponds to observed lower colour temperatures. The effect is not quite so great when high-dispersion spectra are used because occasional, very narrow peaks near the continuum (corresponding to weak absorption) can be found; however, no one claims t h a t a true continuum for a K - t y p e spectrum can be detected in this region. Faraggiana and H a c k also compared line widths in the spectra of these three K - t y p e stars by plotting the central depths, Re, vs. half widths, LJ2, Of the lines. They found t h a t the lines in $ Aurigae are relatively broader t h a n those in the other two stars, and suggested t h a t the continuum of ~ Aurigae is lower t h a n t h a t of the other two stars as a result of broader overlapping wings of the lines, which could be the result of greater macroturbulence or a larger rotational velocity. Victoria 3.2 A/ram spectra of the three stars obtained during total eclipses of 1955, 1961 and 1962 were available and the plot of Rc v s . / [ 2 for fifty selected lines in the region 4180 to 4480 A, is shown in Fig. 12. The data for ~ Aurigae are based on three plates, those for the other stars on only one. I t is seen t h a t the lines in the spectra of $ Aurigae and 31 Cygni are quite comparable and t h a t those in 32 Cygni are appreciably broader. Thus it would appear t h a t the phenomenon noted b y Faraggiana and H a c k is a consequence of the lower-dispersion spectra t h a t were used, and t h a t there is little evidence from high-dispersion spectra for large macroturbulence or rotation in the spectrum of the K - t y p e component of $ Aurigae.* The data given in Table 2 for VV Cephei and e Aurigae must be considered approximate. While additional information concerning the orbit of the secondary star in the VV Cephei system should be gained when the Victoria plates are measured, it is unlikely t h a t the orbit of e Aurigae will be studied in greater detail in the near future, or t h a t additional knowledge concerning the secondary star will be forthcoming. Both systems are much more peculiar t h a n those containing K stars, and it is doubtful whether data concerning * In a recent letter Miss Hack comments: "I think we cannot explain the difference between the widths of the lines of ~ Aurigae and those of 31 and 32 Cygni as a consequence of our lower dispersion because all three stars were studied with the same spectrograph under the same conditions... The lower dispersion could affect the values of c~much more; however those found by us are not very different from yours. Hence I think that ~ Aurigae presents some phenomenon like variable macroturbulence similar to that found for e Aurigae."

K. O. WRmHT !.0o

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179

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components of ~ Aurigae, 31 Cygni and 32 Cygni. their dimensions and masses should be applied to normal stars. The masses quoted for W Cephei depend on the mass ratio and if it is decreased, the masses will a]so be decreased. However, it seems quite probable t h a t the primary star is a supergiant even though Fredrick's (1960) astrometric orbit gives a larger parallax t h a n has usually been considered. The system of e Aurigae remains rather a puzzle since so little information concerning the secondary star is available. With the exception of the Sun and observations of the K-line emission features in latetype stars, most of our knowledge concerning stellar chromospheres comes from observations of ~ Aurigae stars. I t is true t h a t sharp, additional lines, --similar to the chromospheric lines observed near eclipse in the spectra of ~ Aurigae stsrs, - - h a v e been observed in the spectra of some supergiant stars of both early and late type, and in a few variable stars. However, most of these additional lines are displaced relative to the normal wavelength and, in most cases, have been interpreted as arising from distant interstellar material or m a t t e r involving a mass loss from one star to another or to interstellar space; the literature on this subject up to 1962 has been summarized b y W e y m a n n (1963); some additional references arc given b y Wright (1967a). Information concerning the inner

180

The Zeta Aurigae Stars

chromospheres of the ~ Aurigae stars has been reviewed by Wilson (1960), but the only observations t h a t had been published concerning complex K lines were those obtained at the 1951 eclipse of 31 Cygni (McKellar et al., 1959). I t is now known t h a t satellite lines of the C a ~ K line appear near eclipse in the spectra of all three K - t y p e systems and lines of considerable intensity have been seen, for 31 Cygni, at distances of two stellar diameters from the limb of the K star. Some of the lines can be seen for only a few days, but observations made at the 1965 eclipse of 32 Cygni indicate t h a t the lifetimes of others m a y be several weeks. However, it will require intensive observing to learn much more about the duration and the motions of the " c l o u d s " t h a t apparently produce these lines. Since nothing is known about the extent of these clouds, the density of the material producing the lines has also not been determined. Once the radiation from the B star begins to pass through regions of sufficient density to produce absorption, the K-line observations indicate an almost continuous change in the structure of the lines, though no major changes over a few hours' observing have been detected. The opacity of the material in most of these regions is apparently quite small until a few days before totality since the wings of even the central line are absent or very weak. The 1951 observations of the profile of the central line in the spectra of 31 Cygni could not separate it into component lines, but the 1961 observations showed almost conclusively t h a t several components were present. As shown in Fig. 10, the line broadened asymmetrically from November 17 to November 20 and then became narrower on November 23 before becoming broader and finally showing wings on December 18, 4 days before second contact. During the same eclipse, the K line almost disappeared on October 6, when the projected beam from the B star was about one diameter from the limb of the K star. Somewhat similar phenomena occurred during the 1965 ingress for 32 Cygni, as shown in Fig. 4. Thus there is no evidence to indicate major differences in the structure of the atmospheres of these stars and it seems probable t h a t they can be considered representative of giant and supergiant late-type stellar atmospheres. At eclipse it is true t h a t the hot B-type companion in the ~ Aurigae system is only about four astronomical units, from the K star, and therefore the hot radiation from it m a y have some ionizing effects on the outer layers of the giant atmosphere. However, it is unlikely t h a t gravitational effects are sufficient to disturb the dynamical properties of the atmosphere. In the study of the satellite lines observed in the spectrum of 32 Cygni, where the distances are comparable to those in ~ Aurigae, the number of lines showing motions away from the B star are slightly more numerous t h a n those showing motions towards the companion. Therefore it seems probable t h a t conditions found in these atmospheres m a y be considered good first approximations on which theoreticians m a y base their efforts to provide models for late-type supergiant stars. Certainly much theoretical work remains to be done before all the present observations can be explained, but also closer-spaced observations at future eclipses of these stars are essential in order to obtain more definitive data concerning the lifetimes and the structure of the " c l o u d s " t h a t apparently make up a large p a r t of the atmospheres of these giant stars. The author would like to express his thanks to Dr. H. G. Groth who sent the numerous H a m b u r g measurements of K-line intensities obtained at the 1961 eclipse of 31 Cygni, which provided independent confirmation of the trend of the intensity variations. The help of K. H. Hesse in preparing the diagrams and in the computations of stellar diameters, of S. C. Morris in computing the orbital elements, and of S. H. Draper in reproducing the diagrams and spectra is gratefully acknowledged.

K. O. WRIGHT

181

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182

The Zeta Aurigae Stars

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