This dissertation has been microfilmed exactly u received 6 9 -2 2 ,1 7 3

MARKOWITZ, Allan Henry, 1941- A STUDY OF STARS EXHIBITING COM­ POSITE SPECTRA.

The Ohio State University, Ph.D., 1969 A stron om y

University Microfilms, Inc., Ann Arbor, Michigan A STUDY OF STARS EXHIBITING COMPOSITE SPECTRA

DISSERTATION

Presented in Partial Fulfillment of the Requirements for the Degree Doctor of Philosophy in the Graduate School of The Ohio State University

By

Allan Henry Markowitz, A.B., M.Sc.

********

The Ohio S ta te U n iv e rsity 1969

Approved by

UjiIjl- A dviser Department of Astronomy ACKNOWLEDGMENTS

It is a sincere pleasure to thank my adviser, Professor Arne

Slettebak, who originally suggested this problem and whose advice

and encouragement were indispensable throughout the course of the

research. I am also greatly indebted to Professor Philip Keenan

for help in classifying certain late-type spectra and to Professor

Terry Roark for instructing me in the operation of the Perkins

Observatory telescope,

I owe a special debt of gratitude to Dr. Carlos Jaschek of

the La Plata Observatory for his inspiration, advice, and encourage­ ment.

The Lowell Observatory was generous in providing extra

telescope time when the need arose. I wish to particularly thank

Dr. John Hall for this and for his interest. I also gratefully

acknowledge the assistance of the Perkins Observatory staff.

To my wife, Joan, I owe my profound thanks for her devotion

and support during the seemingly unending tenure as a student. I

am deeply grateful to my mother for her eternal confidence and to my in-laws for their encouragement.

ii VITA

October 22, 1941 Born - Jersey City, New Jersey

1963 ...... A.B., University of California, Los Angeles, California

1964-1968 . . . Teaching Assistant, Department of Astronomy, The Ohio State University, Columbus, Ohio

1966 ...... M.Sc., The Ohio State University, Columbus, Ohio

PUBLICATIONS

"Spectrum Variations in the Peculiar A Star 73 Draconis" (co-author with W. K. Bonsack), Publications of the Astronomical Society of the Pacific, Vol. 79, pp. 235-254, 1967.

"New Metallic-Line A Stars", Publications of the Astronomical Society of the Pacific (in press) .

FIELDS OF STUDY

Major Field: Astrophysics Studies in Spectroscopy, Professors Philip C. Keenan and K. Narahari Rao. Studies in Physical Foundations of Astrophysics. Professors Walter K. Bonsack, Eugene R. Capriotti, Wave H. Shaffer, and L. Carlton Brown. Studies in Theoretical Astrophysics. Professor George Collins Studies in Stellar Systems. Professors Carlos Jaschek and Walter E. Mitchell, Jr. Studies in Radio Astronomy. Professors John D. Kraus and H sien C. Ko.

iii TABLE OF CONTENTS

Page

ACKNOWLEDGMENTS i i

VITA i i i

LIST OF TABLES v i

LIST OF ILLUSTRATIONS v i i

C hapter I . INTRODUCTION ...... 1 Definition of a Composite Stellar Spectrum ...... 1 Summary o f P revious W o r k ...... 2 Purpose of This Investigation ...... 7

I I . OBSERVATIONS ...... 8

I I I . SPECTRAL CLASSIFICATION ...... 9 Problems in Classification ...... 9 Description of "Characteristic" Composite Spectrum ...... 10 Classification Criteria ...... 12 Type B Type A Metallic-Line A Stars Type F Type G Type K Type M Comparison with Other Published W ork ...... 16

IV. AXIAL ROTATIONAL V E L O C IT IE S ...... 22 Method of Determination ...... 22 E r r o r s ...... 24 D i s c u s s i o n ...... 25

V. ARTIFICIAL COMPOSITE SPECTRA...... 29 Method of Production and Selection of Types .... 29 Description of Types ...... 34 AOV + GOII AOV + G8III B5V + K3II

iv TABLE OF CONTENTS--Continued

Page

Chapter VI. RESULTS AND DISCUSSION...... 45

APPENDIX A ...... 58

APPENDIX B ...... 75

APPENDIX C ...... 78

REFERENCES...... 80

v LIST OF TABLES

Table Page

1. Mean rotational velocities for e a rly -ty p e s ta r s ...... 26

2. Mean equatorial velocities for early-type stars ...... 28

3. Standards used in making artificial composites ...... 31

U. Relative exposure times for artificial c o m p o s i t e s ...... 33

5. Statistical data for composite s ta rs ...... 55

vi LIST OF ILLUSTRATIONS

Figure Page

1. Comparison of spectral types determined photo- electrically by Bahng with those from this study. Early-type components only ...... 17

2. Comparison of spectral types determined by Bahng with those from this study. Late-type components o n ly ...... 18

3. Comparison of Bahng luminosity classes with those from this study. Late-type components o n ly ...... 19

4. Comparison of spectral types determined by Kuhi with those from this study. Late-type components o n ly ...... 20

5. Comparison of luminosity classes determined by Kuhi with those from this study. Late-type components o n ly ...... 21

6. AOV + GOII urtificial composite spectra ...... 35

7. AOV + G8III artificial composite spectra ...... 36

8. B5V + K3II artificial composite sp e c tra ...... 37

9. Spectral type distribution of early-type c o m p o n e n ts ...... 49

10. Distribution of late-type components with spectral type and luminosity class ...... 51

11. H-R diagram of composite s ta rs ...... 52

vii I. Introduction

Stars exhibiting a composite spectrum form a group intermediate

in separation between close spectroscopic binaries and the relatively

widely separated visual binaries. Despite their potential importance

in leading to a better understanding of the nature of double stars

and their usefulness in providing a check upon current theories of

stellar evolution there has been only one systematic study of com­

posite spectra, performed thirty years ago at this observatory by

J. Allen Hynek (1938),

DEFINITION OF A COMPOSITE STELLAR SPECTRUM

Before Hynek's study and the more recent literature are critically examined, it is necessary to define the term "composite

stellar spectrum". Broadly speaking, the term composite spectrum can

be applied to a spectrum displaying two sets of lines, each ascribable

to a definite spectral type. Within this general definition we may distinguish several classes of systems:

(1) Close binaries consisting of two apparently normal stars

(in the sense that they can be uniquely placed in the H-R diagram) whose spectra cannot be individually observed.

(2) Close binaries in which there are interaction effects be­

tween the components, severely complicating the spectra.

(3) Peculiar objects, including the symbiotic stars, which display a combination of a low temperature absorption spectrum with

high excitation emission lines; 17 Leporis; HZ 9; and others. 2

Among the first class we may distinguish a subclass of stars whose binary nature is evident from the nature of their observed

spectral lines alone, independent of any radial velocity variations.

Also among the f i r s t c la s s we may d is tin g u is h known clo se v is u a l

binaries and some double-line spectroscopic binaries. Others of

the latter subclass can be placed in the second class.

In this study we intend to limit ourselves to those stars

mentioned in the first subclass above. These display a spectrum in which the strength and character of the lines can be attributed to

two different spectral types, independent of any periodic wavelength variations, and which are not known visual binaries. Hereafter, unless otherwise stated, the term '’composite stellar spectrum" will refer to the spectra of this subclass of stars. Close visual binary

systems displaying a composite spectrum are listed separately in

Appendix C but are combined with the program systems for the

statistical analysis in Chapter VI.

SUMMARY OF PREVIOUS WORK

Hynek catalogued all stars known or suspected to be composite,

some 566, from material collected from the Harvard, Yerkes, Victoria,

Lick, Mount Wilson, and Union Observatory files as well as his own material obtained with the Yerkes one-prism, auto-collimating spectro­

graph attached to the 69-inch Perkins reflector located in Delaware,

Ohio. Hynek divided the 566 spectra into nine classes. By far the

la r g e s t number o f th e se , some 371 or 66%, belonged to C lass I . He defined Class I spectra as those "arising from a physical binary whose component stars are not resolvable even with the largest instruments, yet whose relative motion is not great enough to establish definitely the pair as a typical spectroscopic binary." To distinguish these from classical spectroscopic binaries (Class II) and close visual binaries (Class III), they were given the name "spectrum binaries."

That there should be so many of these systems may seem sur­ prising at first but consideration of Kuiper's work on double stars

(1935) shows that this is indeed so. Assuming that the distribution of the logarithm of the semi-major axes of the relative orbits of double stars can be described by a Gaussian distribution, Kuiper showed that these "spectrum binaries", intermediate in semi-major axis between classical spectroscopic binaries and visual binaries, should be as numerous as the former. About one in twenty of the brighter stars can be expected to be a binary system displaying a composite spectrum.

Several factors, however, led Hynek to over-estimate the number of spectra belonging to Class I. (1) A critical reading of the ex­ tensive notes appended to his catalogue reveals that often when there is doubt about the origin of the composite spectrum, especially among the fainter stars (those fainter than Henry Draper visual magnitude

6.0), it has been arbitrarily- assigned to Class I; (2) Many of the

Class I "spectrum binaries" have, since Hynek's work, been shown to be spectroscopic binary systems, thus making the distinction between the two classes somewhat meaningless; (3) Hynek listed as Class I 4 members many o b je c ts now known as m e t a l l i c - l i n e A s t a r s , f i r s t d i s ­ tinguished as a separate group by Titus and Morgan (1940)- These were almost invariably given the spectral classification A + F. A spectroscopic study of the fainter composites in Hynek's catalogue would probably reveal many heretofore unknown metallic-line A stars.

In spite of the above three factors tending to lead to an over-estimate of the number of Class I systems, strong selection effects in both magnitude and spectral class, particularly the latter, are present and can be shown to exclude many valid combinations.

This was most clearly shown in Hynek1s work on artificial composite spectra, also described in his 1938 paper, and confirmed to a limited extent by the author's experience with artificial composite spectra (see Chapter V). These artificial composites indicated two important features. They are: (1) Composite spectra produced artificially, i.e. by superposing the light of two stars in success­ ion on the spectrograph slit, differed in no substantial manner from those found in Nature. (2) Experiments with these artificial com­ posites showed that spectral combinations such as G + K, K + F, F + G, and K + M, are virtually impossible to detect at the dispersion used, and therefore these combinations, if they exist, would be missed in searches for composite spectra. It is interesting to note that combinations such as these are known among wide visual binary systems and

In spite of the many interesting questions raised by Hynek's

study, little work on this subject has been done since then. Bahng

(1958) showed that by multicolor photoelectric photometry, composites

could be classified according to spectral type and luminosity class.

His hope was that photoelectric colors obtained with wide-band filters

would be less sensitive to the veiling and blending effects than is visual

spectral classification based as it is on the relative intensities of

lines and bands. Bahng showed that four colors were necessary to

accurately classify composites; a color in the red region serves as

the best temperature indicator while a color in the ultraviolet is most

sensitive to the luminosity effect. Furthermore, he showed that for

early-type stars the effect of interstellar reddening can be separated

from temperature reddening and the luminosity effect but for late-type

stars the separation cannot be made by the photoelectric colors alone.

It then becomes necessary to rely upon spectral classification for

additional information to decide whether the star is normal or un­

reddened .

Using a 12-inch reflector Bahng obtained the colors of thirty known or suspected composites and assigned two-dimensional spectral

types to 28 of them. It should be noted that a few are known visual

binaries and as such are not included among the program stars of this

study. He determined that 24 of these were composites (including the

close visual binaries) and 4 could be assigned a single two-dimensional

spectral classification. Due to the moderate size of the telescope

this study covered only some of the brighter stars from Hynek's 6

(1938) catalogue. The power of his technique, however, remains clear

in spite of the limited nature of his observing program.

Kuhi (1963) obtained spectra in the visual region (5200 & -

6700 X) a t 88 X/ram with the 7-inch grating spectrograph of the

36-inch Lick refractor. His observing list, chosen mostly from

Hynek (1938), consisted of 64 systems with one member of type B or A

and the other later than F. As in Bahng's case several of his systems are known visual binaries. His MK types refer to the late-type com­

ponent only. In general, they show good agreement with those obtained by Bidelman (1957, 1958) and photoelectrically by Bahng

(1958). In many cases, however, considerable differences were found

from the Henry Draper classifications, upon which much of Hynek*s

statistical discussion was based.

The literature on close binaries exhibiting interaction effects

in their spectra is extensive. Sahade (1960) has suntnarized and dis­ cussed the major theoretical and observational work. Recent work on

the a tte n d a n t e v o lu tio n a ry problem s can be found in Dommanget (1967).

Symbiotic stars are discussed by Struve (1941) and the more modern work is summarized by Sahade (1960). Herbig (1965) has recently discussed the very interesting situation of long-period variables with composite spectra.

Several peculiar stars displaying a composite spectrum have a large body of literature. The interpretation of 17 Leporis as a binary has been discussed by Struve (1932). More recent work by

Slettebak (1950), and with coude< material by Hiding (1966) and Cowley 7

(1967 )j strengthen this interpretation. HZ 9 has been classified as

DA + dMe by Eggen and Greenstein (1965), and is almost certainly a member o f the Hyades c l u s t e r (van A lten a 1969) .

PURPOSE OF THIS INVESTIGATION

The aim of this study is to make a thorough and homogeneous study of composite stars, as defined earlier, complete to a limiting m agnitude. From th i s b a s ic s to re o f d ata se v e ra l q u e s tio n s may be posed which we can hope to answer.

(1) What are the most cotnnon MK type combinations found?

(2) What part do selection effects play in this distribution?

(3) How do the axial rotational velocities of the components compare with those of spectroscopic and visual binaries?

(4) What is the distribution of An among the composites as determined from artificial composite spectra?

(5) What part do selection effects play in this distribution?

(6) What is the relationship of composite stars to spectroscopic and visual binary systems?

An attempt shall be made to answer these questions in the chapters that follow. II. Observations

The observing program was defined to include all known stars exhibiting composite spectra, as previously defined, that are brighter than visual magnitude 7.0, north of declination -25°, and are not known to be visual binaries. This included 74 systems. Spectra were taken of all of these stars, usually one plate per star, with the 8-inch camera of the Y-spectrograph attached to the 72-inch

Perkins reflector of The Ohio State and Ohio Wesleyan Universities at the Lowell Observatory, Flagstaff, Arizona.

This grating spectrograph with the 8-inch camera yields a dispersion of 40 S/mm in the second order. The spectra were taken on

Eastman Kodak IIa-0 spectroscopic plates, presensitized by baking them in an oven for 48 hours at 50° C., and developed for 4,5 minutes in D-19. They were widened to 0.6 mm on the plate in order to in­ crease the accuracy of estimating the spectral types and axial rotational velocities. The projected slit width at the plate was

20^i.

These spectrograms were obtained in three observing runs: the

first during the period 14 February to 25 February, 1968; the second during the period 7 June to 16 June, 1968; the third during the

period 1 October to 10 October, 1968. During the third observing run, artificial composite spectra of the most common spectral-type

combinations were made. The procedures and results are discussed and displayed in Chapter V. 9

III. Spectral Classification

PROBLEMS IN CLASSIFICATION

It has long been realized that the classification of composite spectra is a very difficult problem. Not only are there two different sets of spectral lines present within the same spectrogram but there are interaction effects which must be taken into account. These interaction effects can be placed into two classes: spectral and p h y s ic a l.

The most common example of the first class is veiling. Veiling refers to the weakening of lines of one star by the continuum of the other star. Such veiling noticeably occurs when the early-type component is an early A or a B-type star and may be quite noticeable even when the late-type component is the primary and the magnitude difference is as much as one magnitude. Of course, when the early- type component is the primary the veiling is always very strong and in some cases the lines of the late-type component are nearly invisible,

HD 187949 is an example of this phenomenon. Undoubtedly, many com­ posites have been overlooked in the past, especially on low dispersion plates and in objective prism surveys, because of this. Another important spectral interaction effect is blending. The amount of this effect is directly related to the dispersion of the spectrograph used in the study. At the moderate dispersion used in this study blending was not a severe problem but, for example, it prevented the use of the He I lines 4026 and 4471 in classifying a B-type spectrum 10 when the late-type component was a G, K, or M giant or supergiant.

The physical interaction effects are the interaction effects which are actually occurring to the two components of the system in space. In a few cases there are probably streams of gas being ejected from one or both members forming shells around the system. Other more complicated physical effects may also be occurring. These roost often manifest themselves by the presence of emission components of

Ca II at either H or K or both, or forbidden lines, most notably due to [Fe II]. HD 60414-5, W Cep, and 5 Lac are examples of stars which show these features in their spectra.

Because of these difficulties in classification and because of the unavailability until rather recently of spectra extending into the ultraviolet, most composite stars lacked accurate MK types for both components, especially the early-type. It was not until the spectral surveys of Slettebak (1955), Bidelman (1957, 1958), and, in particular, Kuhi (1963) that some understanding of the distribution of the MK types came about. This study represents the first attempt to determine the MK type of both components of composite stars.

DESCRIPTION OF "CHARACTERISTIC" COMPOSITE SPECTRUM

In much of the discussion below references will be made to a

"characteristic" or "typical" composite spectrum. Before proceeding with a discussion of the spectroscopic criteria used in this study it is necessary to delineate more clearly the above phrase. Upon completion of the spectral classification for this study it became clear that one particular type of spectrum constituted a large fraction of the entire survey. Undoubtedly this is due to selection effects but this point will not be discussed until later.

The spectrum can be characterized as a middle G giant primary with an early A main sequence secondary. The G star is typically one magnitude brighter than the A star in the 4000-5000 X region. The metallic lines of the late-type component are strong and are entirely dominant in the region Hp-He, The Balmer decrement is quite unusual.

The Balmer lines in the ultraviolet are strong but Hf5, Hy, and H6 are quite weak, similar to their appearance in a G giant. In particular,

Hf3 generally appears weak and washed out. He is blended with the H line of Ca II and appears strong because of its natural strength in the spectrum of a G star. The K-line has a definite and unique appearance: a sharp core superposed on a wide, diffuse (variously described as hazy, washed-out, etc.) background. The latter feature is a sufficient but not necessary criterion for compositeness. It represents, however, the single most important criterion in the re­ cognition of a composite spectrum. The spectrum is veiled to a greater or lesser degree depending on the relative strength of the A star spectrum with its continuum (the magnitude difference between the components). Metallic lines can be seen between the higher Balmer lines and on well-exposed plates of composites with G2 and later late- type components, metallic lines can be clearly seen below the Balmer 12 limic. The above description of a characteristic composite spectrum should be kept in mind while reading the remainder of this disser­ tation and in particular the notes on the spectra of the stars listed in Appendix A.

CLASSIFICATION CRITERIA

Because of the nature of the spectra of composite stars it is not always possible to use all of the criteria described by Morgan,

Keenan, and Kellman (1943) and Keenan (1963). Below is a summary of the criteria used to classify the composite spectra.

Type B

The B-type components are nearly always combined with K or

M-type supergiants. The spectrum shows the B star strongly dominating shortward of about H& while the late-type supergiant completely dominates longward of this wavelength. The lines of He I normally play the most important role in the determination of the subdivisions of type B but because of the nature of composite spectra only the

He I line at 3819 X could be used. The absence of a K line core, the strength of the \3819 He I line and occasionally the \4026 line, and the strength and quality of the Balmer lines H8 and above constitute the criteria for spectral sub-class determination among the B-type stars. For the two early B components the luminosity criterion which served was N II \3995/He I \4009. For all the others the appearance of the wings of the higher Balmer lines served as the sole criterion. 13

Type A

The vast majority of early-type components fell within this type.

This, as has been previously mentioned, is undoubtedly the result of strong selection effects since the composite K-line constitutes the principal criterion of classification.

The method of classification is somewhat dependent upon the relative brightness of the A-type component. In general the magnitude difference was not large and both components could be easily seen.

The K-line of Ca II always appears composite with a sharp core super­ posed over a wider, hazy background. The strength and the width of the core primarily determine the spectral subdivision. Care must be taken, however, to allow for the enhancement of the K-line core be­ cause of the presence of the hazy, bright background. The Mg II \4481 line cannot be used since it is always blended with the \4481 line of the late-type component. When the A component is the primary, another criterion can be used throughout the A range: the ratio of

Si II \\4128-30/Mn I W4030-34. This ratio decreases steadily through­ out the A division. In any case the strength of the higher Balmer lines constitute the other important criterion for classification.

They are relatively unaffected by blendings of metallic lines of the late component except for large values of £jn. The luminosity class is determined by the appearance of the higher Balmer lines. In all cases they showed prominent wings and consequently were judged to be main sequence stars. METALLIC-LINE A STARS

The Am stars (Titus and Morgan 1940; Roman, Morgan, and Eggen

1948) were separated out from the composites by the lack of a com­ posite K-line, the apparent noimality of the Balmer decrement, and the clarity of the metallic lines, i.e. the lack of veiling.- The types were determined by comparison with standards and then with well- known Am stars which served as secondary standards. In ten cases,

A + F combinations were found which did not fit the Am criteria.

These were judged to be true composites.

Type F

In ten cases, the F component was the late-type and in five cases it was the early-type component. With the latter, the type was always F2 and earlier. The spectral subdivision was found from the

strengths and width of the K-line core, allowing for the enhancement

by the background, and the strength of the higher Balmer lines. The

luminosity was difficult to judge since the criteria usually quoted;

Fe I X.4171/Ca I X.4226 and Sr I I X.4077/Fe I A.4045 apply to the l a t e -

type component which is generally the primary. With the former case,

the usual criteria may be employed, due allowance being made for any

variations in veiling in the region longward of H5. The major

criterion used was the appearance of the G-band. Luminosity criteria

used were Sr II X4077/Fe I X.4045, Sr II \4077/Fe I X4063, and

Sr II \4077/Ca I \4226. 15

Type G

Numerous metallic lines become visible at this type and some of them prove useful as temperature indicators. The customary MK criteria are the ratio of Fe I X4045/H6, Ca I \4226/Hy, Fe I A.4144/H?) and Fe I X4096/HS. Because of the presence of variable veiling through­ out the spectrum, it was deemed advisable to employ ratios of lines which were closer together. It proved useful to employ the ratios of the strengths of numerous moderately weak lines near HP and Hr to the strengths of Hp and Hr, respectively. The above MK criteria, especially Ca I >4226/Hr, were used as checks. When the G-type was brighter than the early-type component by 0.4 magnitudes or more it was also possible to compare the strengths of the numerous metallic lines between the higher Balmer lines in the ultraviolet to the strength of HP, where the veiling was at a minimum. The luminosity class was determined from the strength of Sr II X4Q77 relative to the

Fe I lin e s AA4045, 4063, and 4071 ; from the r a t i o Sr I I \4 0 7 7 /Ca I

X4226; and from the break in the continuous spectrum at X4215,

Type K

In addition to the criteria mentioned in Class G, the ratio of

Ca I X4226/Fe I X4325 was used. Since the Balmer lines cannot be effectively used as a luminosity criteria in composite stars, the classes were determined from the ratios Sr II x4077/Fe I x.4063,

Sr II \407 7/Fe I \4045, and Sr II \4077/Fe I \4071. 16

Type M

The intensity of the TiO bands provided the principal means of spectral subdivision in this class. The author is especially grate­ ful to Professor Philip Keenan for his assistance in classifying stars in this spectral region.

COMPARISON WITH OTHER PUBLISHED WORK

Only two other previous studies have specifically concentrated on the MK classification of composite spectra. These, as previously mentioned, are due to Bahng (1958) and Kuhi (1963).

In Bahng's study 16 stars were in common with the author's work.

The comparison of the spectral types for the early-type components is shown in Figure 1. Bahng assigned all of these to the main sequence as did the author. A comparison of the spectral types and luminosity classes of the late-type components are shown in Figures 2 and 3.

Although the numbers are small an interesting effect seems present.

Among the early and middle B stars the types assigned by Bahng photoelectricelly differ in no systematic way from those in this study while among the B9-A4 stars, Bating's types are systematically later than those in this study by an average of 3.2 + 0.5 subclass. The late-types show a small scatter with no clear systematic deviations.

A comparison with Kuhi's work for 36 stars in common (Figures 4 and 5) 3hows good agreement for stars later than G5. Earlier than this the agreement is only fair but Kuhi states that his classification is not very reliable due to a lack of suitable lines. IUE . oprsn f ral ye dtr nd oel ri l lly a ic tr c le e to o h p ined determ types l a tr c e p s of Comparison 1. FIGURE SPECTRAL TYPE (BAHNG) A4 A3 A5 AO A2 A6 B3 B9 B7 3 5 7 9 O I 2 3 4 5 A6 A5 A4 A3 A2 AI AO B9 B7 B5 B3 y. ly n o by Bahng w ith those from th is stu d y . E a rly -ty p e components components e p -ty rly a E . y d stu is th from those ith w Bahng by PCRLTP (MARKOWITZ) TYPE SPECTRAL 17 SPECTRAL TYPE (BAHNG) IUE . oprsn f ral ye deemie b Bhg t ith w Bahng by ined eterm d types l a tr c e p s of Comparison 2. FIGURE MO K4 M2 GS KO K2 G6 F 8 G4 GO F 6 6 8 O 2 4 6 8 O 2 4M2 K4 K2 KO G8 G6 G4 G2 GO F8 F6 hs fo this sudy. e-ype cmoet only. components e p -ty te a L . y d stu s i h t from those PCRL YE (MARKOWITZ) TYPE SPECTRAL 18 LUMINOSITY CLASS (BAHNG) Xb XE in FIGURE 3. Comparison of Bahng lu m in o sity c la ss e s w ith ith w s e ss la c sity o in m lu Bahng of Comparison 3. FIGURE 3T y. ly n o those from th i s study. L a te -ty p e components components e p -ty te a L study. s i th from those UIOIY LS (MARKOWITZ) CLASS LUMINOSITY m Xb 19 SPECTRAL TYPE (KUHI) MO KO GO F3 FIGURE 4. Comparison of s p e c tr a l types determ ined by Kuhi Kuhi by ined determ types l a tr c e p s of Comparison 4. FIGURE 4 2 3 5 6 a 9 2 3 4 5 7 4 4 9 2 6 7 8 1 5 1 I

4 6 8 G I 3 5 7 90 4 M 2 I 5M0 4 3 2 1 9K0 8 7 6 5 4 3 2 I 9G0 8 7 6 5 34 F w ith those from th i s stu d y . L a te -ty p e components components e p -ty te a L . y d stu s i th from those ith w y. ly n o PCRL YE (MARKOWITZ) TYPE SPECTRAL 20 LUMINOSITY CLASS (KUHI) Xb 0 1 XX 3T FIGURE 5. Comparison of lu m in o sity c la s s e s determ ined by by ined determ s e s s la c sity o in m lu of Comparison 5. FIGURE Kuhi w ith those from th i s stu d y . L ate-ty p e e p ate-ty L . y d stu s i th from those ith w Kuhi opnns only. components UIOIY LS (MARKOWITZ) LUMINOSITY CLASS nr n b l 21 o l 22

IV. Axial Rotational Velocities

METHOD OF DETERMINATION

Visual estimates were made of the projected axial rotational velocities, the quantity v sin i, for both components of all the composite systems studied. They are listed in Appendix A. These values are the result of three independent estimates made by visual comparison with rotational velocity standards of the same or nearly the same MK type as the composite component. These standards are, w ith few e x c e p tio n s , the same as used for c l a s s i f i c a t i o n purposes as described in Chapter 3. They were all taken by Slettebak with the same equipment as used in this study and their rotational velocities were determined by him in the course of his studies on the variation of rotational velocity throughout the H-R diagram. In most cases the density of the spectra on the program plates was similar to the standard plates used.

An attempt was initially made to obtain microphotometer tracings of the program plates in order to obtain the v sin i from half-width measurements of selected stellar lines. The v sin i could

then be determined by comparison with the half-widths of these same

stellar lines on tracings made from standards of known rotational velocity. This procedure did not prove feasible for several reasons.

First, it was impossible to find lines in the spectra of the com­

posites which could be consistently used throughout the entire study.

Second, variations in the veiling of the spectrum, even between two 23 composites of the same spectral type, yielded widely varying values for the half-widths of the selected spectral lines used, when it could be visually seen that by allowing for the veiling and/or by inter- comparing different parts of the spectrum, the rotational velocities of the two composites were actually similar. Finally, it was realized early that the rotational velocities were in reality rather small or non-existent in most cases and furthermore that the range was rather small. For all these reasons, therefore, it seemed reasonable to obtain the rotational velocity estimates by means of visual comparison only.

In determining the rotational velocities of the composite stars, different lines were used for the early and late-type components, respectively. In the case of the late-type component when the v sin i was not obviously below the plate limit of about 50 km/sec, the

Mg II X.4481 line proved very useful. In addition to this line, the iron lines between and H6, particularly XX4308, 4326, 4384, 4410, and 4415, were used. In most cases, the v sin i was below the plate limit; this is tabulated as 50 in Appendix A.

For the early-type component, the calcium K-line and the Balmer lines, particularly H8 and higher, were the primary lines used to determine the rotational velocity. Despite their broad Stark wings, the Balmer lines can be used for a rough indication of rotation by examining their cores. When the K-line was too faint or absent then the Balmer lines and the He \3820 line if present provided the criteria. 24

It was not possible to use the He lines \4026 and \4471 since they were blended with metallic lines in the late-type spectrum.

In some cases the early-type component is quite veiled due to a rather large Am and there is considerable uncertainty in the value ot v sin i. This is indicated in Appendix A by a value followed by a colon or the word indeterminate.

ERRORS

Before discussing the results, a consideration of the errors would be in order. The errors may be conveniently split into two parts: internal and external. The internal errors are quite small.

The rotational velocities were derived from three independent estimates made at different times and without knowledge of the others.

They seldom differed by more than 15% and in many cases were less.

External errors are a different matter, however. In the first place it is very difficult to estimate the external error since very few of these composites have catalogued values of v sin i. In the few cases which do, the v sin i refers to the late-type component only. In almost every case this is below the plate limit and, therefore, of no real use. Although the method used to determine the rotational velocities suffers from some inherent error, based as it is on visual estimates, in most cases the line widths are small and the resultant v sin i is low, tending to minimize excessive error. 25

DISCUSSION

Table 1 lists the means and standard deviations for groups of early-type stars among composites, visual double stars, and single and non-composite doubles combined. The latter two columns are taken from Meisel (1968) who combined his study with previous studies of

Slettebak and with the tabulated values in the Boyarchuk and Kopylov

(1964) catalogue to obtain the values shown. Only those groups which had significant numbers are listed. In computing the group means and standard deviations composites with v sin i <. 50 were arbitrarily assigned the value v sin i = 50 while those having v sin i < 50 were assigned the value v sin i = 25, In addition, there are four early B-type and four early A-type components. Three of the four early B components have v sin i values that are much smaller than the average for their MK type. One of these, HD 50820, is a Be star and, as a consequence, is probably being viewed at a low inclination angle. The other two either have very small intrinsic rotational velocities or, as is more likely, are being viewed nearly pole-on.

The four early A-type components of HD 157978-9, 59 Ser A,

HD 187949, and £ Cap A have very broad lines. The broad lines are probably due to the presence of a very close companion of the same or nearly the same spectral type as the early-type companion. In three of the cases the presence of a close companion is known. (Batten

1967) and it may be true in the fourth, HD 187949. Therefore, these have not been included in the statistics. The moments of the TABLE 1

MEAN ROTATIONAL VELOCITIES FOR EARLY-TYPE STARS

Composites ^ # < v sin i > Double Stars All Non-Composite Stars Group <7 n < v sin i> o- n < v sin i > a n (km/sec) (km/sec) (km/sec) (km/sec) (km/sec) (km/sec)

B8-A2V 78 +42 39 170 +60 53 160 +60 123

A3-A7V 82 +26 16 150 +50 17 140 +60 49

Results for this column taken from Meisel (1968).

VO O' 27 rotational velocity distribution, as given by Chandrasekhar and

Munch (1950) , are listed in Table 2 for the same groups as in

Table 1.

The results show that except for some B-type components and the four early A-type components mentioned above, all the early-type components in the composite systems studied have projected rotational velocities of less than 130 km/sec and most are between 60 and 110 km/sec. Even allowing for the possibility of large errors in these values, they are still substantially smaller than the means for their respective MK type.

The late-type components, even when they are as early as F-type giants, have little or no rotation. This is contrary to the situation in single or wide double star systems and further confirms the in­ trinsically slow rotational nature of composite systems.

Although these systems appear to be slowly rotating most of them have not reached a state of axial-orbital synchronism as has been shown to exist in some spectroscopic binary systems and more recently in eclipsing binary systems (Olson 1968). In general the periods, when known, are quite long, averaging about six years for the program systems. The projected rotational velocities would have to be considerably less than the tabulated values in order for the systems to reach an axial-orbital synchronous state. TABLE 2

MEAN EQUATORIAL VELOCITIES FOR EARLY-TYPE STARS

Composites All Non-Composite Stars a cr Group < V > ^ V > v V > < V > (km/sec) (km/sec (km/sec) (km/sec)

B8-A2 V 99 +44 200 +50

A3-A7 V 104 +15 180 +60

Taken from Meisel (1968),

N) 00 29

V. A rtificial Composite Spectra

In order to obtain another parameter of information from the program plates, i.e. the magnitude difference between the components, artificial composite spectra were constructed. They were also used as an aid in determining the MK types of the program stars. These artificial composites were made during the October 1968 observing run using the same equipment, plates, method of exposure and de­ velopment as with the program composites, so as to insure complete homogeneity throughout.

METHOD OF PRODUCTION AND SELECTION OF TYPES

The artificial composites were constructed by allowing the light from an MK standard to fall on the spectrograph slit and trail­ ing the star along the slit for a predetermined length of time, then ending the exposure by closing the camera shutter and (in order to minimize seeing differences) quickly moving the telescope to a nearby

MK standard, re-opening the shutter, and trailing this star along the slit for another predetermined length of time. The iron arc com­ parison spectra were impressed on the plate before the first exposure and after the second exposure. With practice the entire operation of moving the telescope from one star to another required no more than two minutes and often was accomplished in one minute. The artificial composites were taken only on nights of steady seeing and good trans­ parency and only when both the MK standards were near the zenith. 30

Initially it was hoped that wide visual binaries with well-

determined MK types and known rotational velocities could be used as

the standards for constructing the artificial composites but this did

not prove feasible. Using these wide binaries would nearly eliminate

any unwanted factors such as varying seeing in different parts of the

sky and over the period of time necessary to move the telescope from

one component to the other. There were not, however, wide binaries with the proper MK types and rotational velocities in the accessible

part of the sky during the observing period. Instead, it was necessary

to use MK standards or other stars with well-determined MK types and

known rotational velocities which were accessible and positionally

as close to each other as possible.

By the time the observing program for the artificial composites

had been worked out enough of the MK classification had been done to

let the writer know the most common combinations of MK spectral type

occurring among the program stars as well as to realize that most of

the early-type components were rotating only slowly. Therefore, the

three most common MK combinations were chosen as representative of the

group and standards were selected on this basis. The data on these

standards are listed below. TABLE 3

STANDARDS USED IN MAKING ARTIFICIAL COMPOSITES

STAR a(1968) &(1968) MK TYPE v sin i VB (km/sec)

p Aur 5h 19?5 +4l°46' B5 Va 90 :a 5.093 4,91a

29 Vul 20 37.0 +21 05 AO Vb 48b 4.82b 4.80b u LO Q Sge 19 38.5 +17 57 GO I I 3 V 4.38b 5,15b

24 Vul 20 15.5 +24 34 G8 I I I C < 20 :C 5 .3f* 6.43C

( Aur 4 55.0 +33 06 K3 I ! 3 < 15C 2,67b 4.21b

The sources are as follows: SPECTRAL TYPE MAGNITUDES a. Johnson and Morgan (1953) a. Bright Star Catalogue (Hoffleit 1964) b. S lettebak (1954) b. Arizona-Tonantzintla Catalogue c. Roman (1952) (Iriarte et al. 1965) ROTATIONAL VELOCITY c. Harvard Photographic magnitude a. Slettebak and Howard (1955) transformed to blue magnitude. b. S lettebak (1954) c. Herbig and Spaulding (1955) 32

A rtificial composites were constructed using the above standards for the following MK combinations: AOV + GOII,

AOV + G8III, and B5V + K3I1. For each combination a set of seven plates was made, each with a different chosen magnitude difference between the components. These magnitude differences were: 0.0, + 0.5,

+ 1.0, and + 2.0. The relative exposure times given to the standards,

and T 2 _> were calculated from the formula

T. (d-d ) + (B -B ) = (2.512) l 2 where t*ie desired difference in photographic magnitudes, in the artificial composite, and (B^-Bj) is the actual difference in blue magnitudes taken from Table A. The calculated relative exposure times are listed below.

In all 25 artificial composite spectra were made, 21 of them using the MK type combinations described above and A others of varying combinations and lesser quality to test the feasibility of detecting other combinations. Every precaution was taken to insure that the procedures outlined above were rigorously followed so that any differential seeing and transparency effects were kept to an absolute minimum and that the total exposure times were such that the density of the artificial composite plates closely matched those of the program plates with which they were to be compared. The values

listed in column 7 of Appendix A were determined by visually com­ paring each program composite with the most appropriate set of TABLE 4

RELATIVE EXPOSURE TIMES FOR ARTIFICIAL COMPOSITES

STAR MK TYPE

29 Vul AOV 0.0 0.5 -0.5 1.0 -1.0 2.0 -2.0 (VL. di>1 (B1-B2) = -0.35 or Sge 0.72 1.15 0.42 1.82 0.29 4.60 0.11 GOII Tl /T2

29 Vul AOV (d2-d1) 0.0 0.5 -0.5 1.0 -1.0 2.0 -2.0 (b1“b2) = ‘ l *63 24 Vul G8III 0.22 0.35 0.14 0.56 0.09 1.41 0.04 V T2

p Aur B5V (B ^ ) = 0.70 0.0 0.5 -0.5 1.0 -1.0 2.0 -2.0

I Aur K3II 1.90 3.00 1.20 4.80 0.76 12 0.29 Tl /T2 34 artificial composite stellar spectra.

The artificial composites are displayed in Figures 6, 7, and

8. The paragraphs which follow give a description of the salient features of the spectra as they appear on spectrograms with a dis­ persion of 40 S/tnm and the manner in which they change as the magnitude difference is varied. A positive (+) Am indicates that the early-type component is the primary while a negative (-) Am indicates that the late-type component is the primary.

DESCRIPTION OF TYPES

AOV + GOII

The combination of an early-type main sequence component near

AO with an early G-type giant component is the most common combination

found among the system s in th is survey.

Am = o

The spectrum has the characteristic composite appearance as described in Chapter 3. The K-line has a sharp core superposed on a

broad hazy background. The spectrum in general is veiled, though not

severely. The HP, Hy, and H£> lines resemble those in a GO II spectrum

but of course their contrast is somewhat diminished. He is about one-

third less broad than in the spectrum of the GC II standard and the

sharp H line of Ca II is blended with its core. Longward of HC, the

metallic spectrum matches in every respect that of a GO II star except

that the contrast is down by one-half due to the veiling. Shortward

of H£, however, the Balmer lines are much stronger than in a G-type 35

O O m L^i cm o o CM < I + + II II II II ii B B I <\

FIGURE 6. AOV + GO II A rtificial Composite Spectra. 36

0 o l / ' l o o tM ■—! o O ^—1 CM i I + + II II n II It 1 I % 1 1 1

9 I I I I * ■ « 1 • l

FIGURE 7. AOV + G8III A rtificial Composite Spectra, 37

o o LO O

CM p 1 O O O - * » 1 i

II 11 II II II II

Am 1 1 1 1 s

FIGURE 8. B5V + K3II A rtificial Composite Spectra. 38 spectrum though they are not nearly as strong as in an AO V spectrum.

This is due to the superposition of numerous metallic lines fiom the late-type spectrum. In this case the main classification criterion for the early-type component is the strength and width of the K-line core. The position of the Balmer limit is also uncertain, due to the blending of the higher Balmer lines with numerous metallic lines from the G spectrum.

/an - -0 .5

The spectrum is very similar to the previous case. The main differences are: the veiling is only about one-half that of the

= o case; the K-line background is slightly wider and stronger; the core is unchanged ; and the Balmer lines shortward of H€ are slightly sharper.

Am = -1 .0

Only a slight amount of veiling is present. The Sr II X.4077 line seems to stand out more clearly against its background than in previous cases. The K-line background is not as diffuse as in

Am = -0.5; its edges are becoming defined. The core is unchanged.

The metallic lines between the higher Balmer lines are very definite even when they fall in the wings of the latter. The strength of the higher Balmer lines over those in a G-type spectrum is still very pronounced, however. 39

Am = -2.0

The strength of the K line background is about twice that in the A m = -1.0 case and the line has quite definite boundaries. The core is still visible but its brightness is very nearly the same as that of its background. No trace of veiling remains in the spectrum.

The higher Balmer lines very nearly resemble those in a G giant spectrum; the only definite classification criterion remaining for the early-type component is the width of the K-line core. At lower dispersions the presence of an early-type component could easily be overlooked.

A m = +0.5

The K-line background is only faintly visible; the core is pre-eminent. Hp, Hy, and H6 are still "composite" in appearance though Hp is noticeably stronger than in Aan = o. The higher Balmer lines, shortward of the K-line, are stronger, sharper, and in particular less blended than for the case Am = o. The metallic spectrum of the G star is rather strongly veiled, though most of the lines are still visible; the most pronounced features are \\4481,

4383, 4325, the G band, 4226, 4215, 4077, 4045, some faint lines blended with the violet side of H8, 3630, and 3617.

Am = +1.0

No trace of the K-line background remains. The Balmer lines are nearly as strong as in an AOV standard: only H£ and Hy still show some blending. The G spectrum is very weak. The G-band is 40 weak and is scarcely visible* The most pronounced features are:

\\4481 (most of the strength is contributed from the A-star), 4326,

4383, 4226, 4215, 4077, and 40^5.

4jn = + 2 .0

Blendings of the early Balmer lines, HP, Hy, and Hb, are only barely seen* The most obvious is Hy. The higher Balmer lines practically match those in an AOV standard. The metallic line spectrum of the G star is very weak. The lines most easily visible are: \\4481, 4383, 4326, 4226, 4215, 4077 and 4045. There is virtually no trace of the G star in the spectrum short of H€.

AOV + G 8III

£jn = o

The spectrum, as in the analogous case of the AOV + G0II com­ bination, shows the composite character. The K-line has a sharp core superposed on a broad, hazy background. Veiling of the spectrum is quite noticeable. Hp and Hy are about the same strength and character as in a G8III standard; H6 is a bit stronger. He is about one-third

less broad than in the spectrum of the G8III standard. The higher

Balmer lines are much stronger, and more e^gj.ly visible because of

the relative lack of blending with metallic lines in this spectral

region* The metallic line spectrum has only about one-third the

contrast of the G8III standard, although most of the lines are still

easily seen. 41

Am - -0.5

The K-line core is slightly stronger probably due to the greater strength of the background, though the latter change is not clearly seen. The Balmer lines are very nearly unchanged from the

Am = o case except for He which is nearly as broad as in the G8III standard. The major difference is the contrast and visibility of the G spectrum. It is now about one-half the contrast of the G8III standard. The metallic lines are still almost invisible shortward of

4045 X. The CN band a t 3883 X is clearly seen.

Am = -I.Q

The K-line core although still visible, is much less obvious than in Am = -0.5 and the hazy background is stronger and has more definite boundaries. The Balmer lines HP, Hf, and H& are very similar to the G8III standard though H6 looks a bit stronger. Some metallic lines are now seen in the higher Balmer line region making the hydrogen lines less easily distinguishable. The contrast of the metallic lines in the 4045 X - 5000 X region is down about one-fourth from the G8III standard. The CN band at 3883 X is quite prominent as are numerous metallic lines shortward of the Balmer limit.

Am = -2 .0

The spectrum is similar to an ordinary G8III standard. The major differences are: The K-line is not quite as wide and strong and there is a hint of a core seen under high magnification; the metallic line contrast is not quite as sharp; and the Balmer lines 42 are a little stronger and more well-defined.

Am = +0.3

The K-line core is sharper and more definite than in the

Am = o case; the hazy background is very slightly fainter. The rest of the spectrum closely resembles that of Am = o except for a small increase in veiling of the metallic spectrum of the G star.

Am J +1.0

The K-line background is barely visible. The Balmer lines, Hp,

Hy* and H& are still noticeably weaker than in an AOV standard but the higher members of the series are nearly as strong as in the standard.

The metallic line spectrum is heavily veiled but several lines are nevertheless seen. They include: \\4481, 4384, 4326, 4226, 4215,

4077, and 4045. In addition several metallic lines between HP and

5000 X as well as \3630 and \3617 are easily visible due to a lessening of the veiling in these regions.

= + 2.0

The spectrum clearly resembles that of a AOV standard. The only indications of the G-type star are: a slight weakening of HP and H&, barely visible X.X.4326, 4386, 4045, 3630, and 3617 lines.

B5V + K3II

This combination is representative of another common combination found among the program composites: that of a middle B main sequence and a high luminosity late-type companion. 43

Am = o

The spectrum can be neatly divided into two parts. Shortward of 4000 X it matches a B5V standard while longward of this it matches a K3II standard. Close examination shows some enhancements of the features at \402fe and \4471 in the K spectrum due to the presence of strong He I lines from the B spectrum.

Am = -0 .5

The K-type spectrum can be seen down to the K-line; shortward of that it is still nearly invisible except for some faint lines beyond the Balmer lim it.

Am = -1.0

Metallic lines of the K star are seen between all the higher

Balmer lines, though the lines themselves are still strong and de­ finite. The K-line is very light but its boundaries are more definite than in the Am = -0.5 case.

Am = -2 .0

The K spectrum is dominant throughout the entire spectrum but the higher Balmer lines can still be made out. The K-line has clear boundaries but is slightly weaker than in a K31I standard. These two criteria remain the only way to recognize the presence of the B-type spectrum .

yan = +0.5

The K spectrum is clearly seen longward of Hy and seen through veiling to about 4045 X. Shortward of that point the B spectrum is strong and pure with He I \4026 clearly seen.

Am = +1.0

The entire spectrum is veiled to a greater or lesser extent

Most of the K spectrum longward of H5 is still seen; in particular

\4471 is greatly enhanced over that in a K3II standard.

Am = +2.0

The B spectrum is extremely dominant. The K spectrum can only be seen, through heavy veiling, longward of Hy. All the He I lines in the B spectrum are clearly seen including \\3820, 4009,

4026, 4X44, 4387, and 4471. 45

VI. Results and Discussion

The results of this study are tabulated in Appendix A. Therein is contained a list of all the stars known to the writer which have been reported to exhibit a composite spectrum as defined in the

Introduction and fall within the position and magnitude limits of the study. Hynek's (1938) catalogue served as the basis of the list but many other systems were culled from numerous other sources in the literature.

The successive columns list: (1) the catalogue number by which the star may be identified in the notes which follow; (2) the

Henry Draper Catalogue number; (3) the Bright Star Catalogue number;

(4) the star name^ if any; (5) the spectral types of the primary and secondary, respectively, as assigned by the writer; (6) the pro­ jected rotational velocity, v sin i, of the primary and secondary as assigned by the writer; (7) the difference in photographic magnitudes between the two components of the composite in the sense secondary minus primary, as assigned by the writer using the artificial com­ posite spectral standards as described in Chapter V; (8) and (9) previously determined MK types with the literature references; (10) and (11) previously determined rotational velocities, if any, with the literature references.

Following the table are notes on the spectrum of each star.

The notes give a general description of the spectrum of each star with particular emphasis on the appearance of the calcium K line. 46

In addition to this, mention is made of any unusual spectral features such as emission lines or enhanced absorption lines. Also included is an indication of the spectroscopic binary nature of the star

(denoted SB). The notation is either followed by a number in parenthesis, indicating the radial velocity variation as recorded in the General Catalogue of Stellar Radial Velocities (Wilson 1953), or the period as listed by Batten (1967). Finally, there is in some cases a reference to the double star measurements of Raymond Wilson,

Jr. Wilson (1941, 1950, 1951, 1952, 1954, 1955) in a series of papers reported measurements on close double stars made primarily with an interferometer attached to the 18-inch Flower refractor of the University of Pennsylvania. Included among these binaries were a few stars taken from Hynek's (1938) catalogue of composite stars.

Many of these were found to be single by Wilson but on some occasions he reported finding small separations, or at least upper limits to the separation, between the components. Unfortunately, he often lists the observation as of poor quality or indicates widely differ­ ing position angles for two observations taken in succession. His results have never been confirmed by anyone else and are the only entries for these stars in the Index Catalogue of Double Stars

(Jeffers et al., 1963). For these reasons it was felt that these stars should be kept in the writer's program list although probably some time in the near future some of these stars will be shown conclusively to be binary. The best candidates seem to be 31 Cyg A,

9 Cyg, r Per, and perhaps 47 Cyg. 47

Appendix B gives a list of all stars falling within the position and magnitude limits of this study which have been reported to have a composite spectrum but upon subsequent analysis have been shown either to be otherwise or to be close visual binaries. Many of these stars are taken from Hynek's (1938) catalogue and have long been known to be other than composites, according to the definition given in the introduction to this study.

The successive columns in both tables list: (1) the catalogue number by which the star may be identified in the notes which follow; (2) the Henry Draper Catalogue number; (3) the Bright

Star Catalogue number, if any; (4) the star name, if any; (5) the spectral type assigned by various recent observers including the writer. The notes give the references for the classification, in­ formation on the magnitude differences and separations of the visual binaries, and the various types of the metallic line A stars. Unless otherwise stated, the metallic line A types refer to those assigned from the K-line, the Balmer lines, and the metallic lines, re­ spectively .

It is interesting to examine the distribution of objects listed

in the above table of suspected composites. Of the 22 stars, 11 are metallic line A stars, two of which were found in this survey (numbers

10 and 15) two others are peculiar A or B stars, four are high

luminosity stars mistakenly classified as composite, three are

apparently single stars which may be assigned a unique MK type, and

two are rapidly rotating shell stars. 48

It is interesting to note that the recently reported Am stars

(Markowitz 1969} show the scandium anomaly discovered by Conti (1965), although the case of 17 Hyd B is somewhat uncertain since the Sr II

\4215 line is slightly marred by a plate defect. It is somewhat surprising that the scandium weakness can be detected at this dis­ persion but this effect was checked by comparing the plates against early A standards as well as "classical'1 Am spectra used as secondary classification standards.

Figures 9 and 10 show the results of the spectral classi­ fication. The first figure is a histogram showing the distribution of the spectral types of the early-type components. The luminosity class was always found to be V. Although there is difficulty in de­ termining the luminosity class of some of the early-type components, particularly when the early-type is the secondary and the Am is greater than 0.5, the H-R diagram in Figure 11 seems to confirm the above finding. This is discussed further below.

It can be seen from the distribution that the vast majority of the early-type components lie within the range B9 through A4 with peaks at AO and A2. This confirms earlier findings and can be shown to be largely the result of selection effects since the presence of a composite K-line with its sharp core and the Balmer decrement, as previously der.r ribed, in the spectrum of a later-type giant or super­ giant are the principal criteria of compositeness. These two criteria are most strongly manifested in early A-type dwarfs. NUMBER 12 15 IU E . pcrl ye sti i f al-ye components. early-type of n tio u trib is d type Spectral FIGURE 9. 0 3 6 9 O 2 4 6 8 O 2 4 6 8 O F2 FO A8 A6 A4 A2 AO B8 B6 B4 B2 BO SPECTRAL TYPE ID 50 Figure 10 shows the distribution in spectral type and luminosity class of the late-type components. It shows that the majority of components are giants falling within the range G0-G8 with a peak at

G5. It is very similar and is to be compared with Figure 1 of Kuhi

(1963). Although his was not a complete study down to a limiting magnitude his distribution is very similar with the exception that no late-type dwarfs were found in the present study. It is to be noted that Kuhi states that his classification is not very reliable for types earlier than GO and this is where all but one of the dwarfs he finds fell.

A Hertzsprung-Russe11 diagram showing the positions of both components of all the composite stars studied is shown in Figure 11.

The Roman numerals lb through V are the sequences of supergiants through main sequence, respectively, from Keenan (1963), Although the difficulties of assigning accurate luminosity classes to A-type spectra are well-known, especially when they are further complicated by a composite spectrum, the early-type components were all judged to be dwarfs and placed on the Keenan main sequence in Figure 11. The

AM = Am = AV was found from the formula AV = AB - A(B-V) where the v v AB is the difference in blue magnitudes and A(B-V) is the difference in color index between the components. The AB used was the found by means of the artificial composite spectra as described in Chapter V.

The B-V values for the components were taken from Johnson (1963).

Although the errors in AM^ by the nature of its determination preclude the possibility of an independent check of the Keenan giant and supergiant sequence calibration, they are not so large as to mask

the rather clear result: In every case that a component was classified LUMINOSITY CLASS TS' Ib- 1 F3 F3 FIGURE 10. ---- 456789 60456789 9 8 IKO 7 6 5 4 3 2 I 1 ----- 1 ----- ah a rpeet oe ar ^ r. ta s one represents bar Each D istribution of late-type components components late-type of istribution D 1 ----- 1 ----- 1 ----- 1 ----- 1 ----- 1 ----- 1 ------SPECTRAL TYPE I 1 I ! with with ------pcrl ye n lmnst class. luminosity and type spectral 1 ------1 ------1 —I - 1— 2345 -- 1 ------1 - 1 ------MO I 1 - 1 - 1 ---- 1 ----- 234 1 ----- 1 ----- 1 LUMINOSITY CLASS - w b- lb - L X - K F3 F3 FIGURE 10. D istrib u tio n of late-type components with sp ectral type and luminosity c lass. lass. c luminosity and type ectral sp with components late-type of n tio u istrib D FIGURE 10. 456789 GO456789 I 9 8 KO 7 6 5 4 3 2 12345 MO I 234 ah a rpeet oe t r. sta one represents bar Each PCRL TYPESPECTRAL 52

H -R DIAGRAM OF COI

- 2.0

BO Bl 6.O1

.FIGURE 11. H-R Diagram of Composite Star Components, Early-type components have been fit to Keenan main sequence (see te x t) . MPOSITE STAR COMPONENTS 53 by visual means as a giant, for example, the Am transformed to a Pg places it within the giant range. The same is true for luminosity classes IV, II and lb. In only a few cases is there any ambiguity and then it is when a component was assigned a luminosity class of

III-IV and it seems to fall within a few tenths of a magnitude of the III line.

Figure 11 is similar in form to H-R diagrams of visual binary components (Bidelman 1958; Stephenson 1960; Slettebak 1963) and appears to confirm current ideas of stellar evolution. The primary components, being more massive, have evolved off the main sequence and become the late-type giants and supergiants now observed. As

in the case of visual binaries, the evolution of the components in

the systems studied seems to have proceeded independently of one another, except as noted below. This is entirely understandable

since the separation of the components of the composite systems is

large in comparison with the maximum separations necessary for one

component to evolve to its Roche limit, overflow, and significantly

effect the evolution of its companion (Plavec 1967). The spectra,

in general, confirm this by displaying a lack of unusual spectral

features associated with semi-detached and contact binary systems.

One interesting feature of the H-R diagram of the components

is the relatively large number of stars that appear to lie in the

Hertzsprung gap. This confirms an earlier finding of Kuhi (1963)

and is in marked contrast to visual binaries, as he showed. This

seems to indicate that, in some cases, the present late-type com­ ponents of the composite systems have not evolved as far off the main sequence as they have in visual binary systems. This suggests that

the closeness of the components may have affected their evolution,

albeit to a far lesser extent than in close spectroscopic binaries.

The fact that the early-type composite components rotate more slowly

than their corresponding types among visual binary and field stars

(see Table 1) suggests that this may be the case.

In order to understand better the nature of composite stars,

Table 5 has been compiled. It is similar in some respects to

Table 3 of Hynek (1938). The columns show the respective number of

various classes falling into the magnitude intervals indicated. The

first column lists the number of composite stars that fall within

the position limit of this study and which are not known visual

binaries. The second column shows the number of close visual binaries

exhibiting composite spectra within the position limit of this study.

The third column shows the number of systems in the first two columns

which have an A-type component. The last column shows the total

number of all stars in the sky within the position limit of the

study. These data were taken from Haramundanis (1967) which contains

a statistical discussion of the Smithsonian Astrophysical Observatory

Star Catalog (1966). TABLE 5

STATISTICAL DATA FOR COMPOSITE STARS

Class No* of No. of No. of Vis. Composites Close Visual A- Type No, of Mag. Binaries Components A ll Stars

- 1.0 0 1 0 9

1.0 - 2.0 0 0 0 16

2.0 - 3.0 2 0 1 72

3.0 - 4.0 9 0 5 205

4,0 - 5.0 11 1 6 688

5.0 - 6.0 19 8 17 2240

6.0 - 7.0 33 8 36 6930

Total: 74 18 65 10160

Ui u* 56

An analysis of the data presented reveals several interesting features. From the observed ratios of the number of composite stars

(including close visual binaries displaying composite spectra) to the total number of all stars, among the brightest magnitude intervals, we should expect to see at least 356 composites brighter than 7th magnitude; instead we find 92 of them. Thus, we should expect to find at least four times as many composites brighter than 7th magnitude than are presently known. From studies of artificial com­ posite spectra it seems fair to say that the sample of composite stars is limited by marked selection effects in magnitude and spectral type.

The extent of the magnitude selection effect is difficult to judge but it seems that it is not as severe as the spectral type selection effect. With favorable combinations, e.g. a B-type combined with a late-type giant or supergiant, differences of up to three or more photographic magnitudes (six visual magnitudes) may be present and both types can still be discerned on the plate. Even with the more commonly found combinations, differences up to three visual magnitudes may be present.

The extent of the spectral type selection effect can be seen by noting from Table 5 that slightly more than two-thirds of all the composites (including close visual binaries) have an A-type com­ ponent. To be sure this is not as large as the over 907, found by

Hynek, due in large part to his inclusion of metallic-line and

peculiar A stars, but it nevertheless demonstrates the fact that 57 composites with an A-type component are much easier to find, especi­ ally on objective prism plates. Among the composites brighter than the fifth magnitude the fraction with an A-type component is one-half.

Fainter than this the fraction grows steeply.

We may calculate the mean period of the 24 composites which are spectroscopic binaries, and from this, assuming typical masses, the mean separation. The mean period is 2208 days or 6.05 years. For typical masses the mean separation lies between 5.8 and 9.0 A.U.

We may summarize the major conclusions of this study as follows:

(1) The spectrum of a composite star typically consists of a

G-type giant combined with an early A-type dwarf. The spectral type distribution is strongly affected by selection effects in magnitude and spectral class. There are estimated to be at least four times

as many composites brighter than seventh magnitude than are presently

known.

(2) The mean rotational velocities of the early-type components

are substantially less than the means for their corresponding MK types

among wide visual binaries and field stars.

(3) A significant number of composite components populate the

Hertzsprung gap. APPENDIX A: LIST OF STARS WITH COMPOSITE SPECTRA BRIGHTER THAN mv = 7.0 AND NORTH OF -25° DECLINATION

Previously Previously Determined MK TYPE v s in i (km/sec) Determined v sin i NAME PRIMARY SEC. PRIMARY SEC. Am. MK TYPE REFERENCES REFERENCES NO. HD HR ph (km/sec)

1 1142 G8UI F2V: £ 50 ind. 1.3: C8III+F Bidelman (1957)

2 1952 F 6III A4V 70 100 0 .1 ::

3 4615-6 C8III A2V 90 125 0.7

4 4775-6 233 B9.5V GOIII-IV £ 5 0 < 50 0.3 COIII+A4V Bahng (1958) 0 Boyarchuk & c o m Kuhi (1963) Kopylov (1964)

5 5621 A8V GOII-III 120 70 0.8

6 9352-3 439 K 3lb-II B9V £ 5 0 200 0.6 K3lab Kuhi (1963)

7 13474-5 640 55 Cas B9V GOII-III 85 70 0.2

8 14262-3 676 A1V: F 3III £ 50 110 0.1 F6IV Kuhi (1963)

9 17245-6 A1V G2III £ 50 80 0.4

10 17878-9 854 t Per G 5III A4V £ 50 £ 50 1.3 C5III+A Slettebak (1955) < 25 S letteb ak ( 1955) G4III+A4V Bahng (1958) < 25 Boyarchuk & Kopylov (1964) G5IV Kuhi (1963)

11 18925-6 915 I Per G 5III A2V < 50 4 50 1.0 A5+G5III Slettebak (1955) o 25 S lettebak (1955) G8III:+A3 Stebbins 6 0 Boyarchuk 6 Kron (1956) Kopylov (1964) C8III+A3V Baling (1958) G 8III Kuhi (1963)

12 19926-7 958 K lIIIcp A6V £ 50 120; 0 .9 K SlIIep Kuhi (1963)

13 20084 965 COll FO: V SO £ 50 2 .5: G 8 IM II Roman (1955) G 3p:II Keenan 6 R oller (1953) I

Al'l'KNMX A: Continued 58a

Previously Previously UK TYPE v sin i (km/sec' Determined Determined PRIMARY SEC, PRIMARY SEC. Ani KK TYPE REFERENCES v sin i REFERENCES ______(km/sec)

14 23089-90 1129 B9V G211-I1I i 50 £ 50 0.3 A1-K12IJ] Slettebak < 25 Slettebak (1955) (1955) G0111+A3V ^‘ ' „ (1958) Boyarchuk 6 Kopylov (1964 GOIII Kuhi (1963)

15 23838 1176 G2II1 F2:V i 50 85: 0.8: G2I117+A? Appenzeller (1967)

16 25555-6 1252 36 Tau A B9V G11II 160 £ 50 0.3 GOIII+A4V Bahng (1958)

1? 26673-6 1306 52 Per G51I A2V i 50 100: 1.7 A2V+G51b S lettebak < 25 Slettebak (1955) (1955) G5II+A,B Bidelraan Boyarchuk 6 (1957) Kopylov (1964) K2III+A6V Bahng (1958) G8III Kuhi (1963)

18 29094-5 1454 58 Per K O II-III B9V i 50 £ 50 1.2 G8II+B Bidelman < 25 Slettebak (1954) ( 1955) A+GSlb-11 S lettebak < 25 Boyarchuk 6 (1955) Kopylov (1964) K4III+A3V Bahng (1958) KOI 11 Kuhi (1963)

19 32068-9 1612 [ Aur K5II B5V £ 5 0 3.1 K5II+B Stebbins 6 41 Boyarchuk 6 Kron (1956) Kopylov (1964) K4II+B8V (1958)

20 34318-9 G81II AOV £ 5 0 GBIH Kuhi (1963)

21 39118-9 2024 G8III AIV £ 5 0 85 0.7 G8IV Kuhi (1963)

22 39286 2030 B8V GZHIe 230 B9V +K? Osawa (1959) 350: Palmer et a l. (1968) GSIIIe Kuhi (1963) Palmer et al, (1968)

23 47579-80 G5III A3V 70 Kuhi (1963)

24 48953-4 G5Ia A5V £ 50 » 2 I

59 APPENDIX A: Continued

Previously Previously MK TYPE v sin i (km/sec) Determined Determined PRIMARY SEC. PRIMARY SEC. to UK TYPE HR ph REFERENCES v sin i REFERENCES (km/sec)

25 50820 2577 B3eV 130 ind, 2 ;; B3eV+K2lI U n d e r h i l l (1954) B3;p Mendoza (1958)

26 52690 MOIb AOV: i 50 ind. 3: Mllb+A.B Bidelman ( 1 9 5 7)

27 52822-3 G8III AlV < 50 < 50 0,9

28 60414-5 2902 M2Iabep B ind. ind. M2Iabep+B Bidelman (1954) M2lbep Kuhi (1963) Mllabepi B2 C.&M.Jaschek (1963)

29 69479-80 G5III A2V: < 50 ind, 0.3 G5III Kuhi (1963)

30 70442-3 3279 G8III AQ:V < 50 80: 1.5: G8III Kuhi (1963)

31 74228-9 3450 A3V GOIII 130 70 0.6 F8V Kuhi (1963)

32 76370 3553 G5III AOV <50 < 50 0.6 Am(A3-A3-F2) S lettebak < 25 S letteb ak (1963) (1963)

33 83808-9 3852 o Leo AlVAlV F6II < 5 0 < 5 0 0,2: A2+F6II-III S lettebak < 25 S lettebak (1955) (1955) F8III+A5V Baling (1958) 15 Boyarchuk & Kopylov (1964) F6III Kuhi (1963)

34 102509 4527 93 Leo GG 4IIM 4IIM V V A7V < 50 110 0,8 A+G51IHV S lettebak < 25 Slettebak (1955) (1955)

35 80 4707 12 com A2 A2v V COIII-IV 80 < 50 0.1 A3VIC0IIMV Herbig A < 25 S lettebak Turner (1953) (1955) A5+G51II S lettebak 20 Boyarchuk A ( 1965) Kopylov (1964) c o m Kuhi (1963)

36 108464-5 A7V com 60 £ 50 0.1 A i m 'l l )X A: Continued 60

P reviously P reviously MK TYPE v sin i (km/sec) Determined Determined PRIMARY SKC. PRIMARY SEC. Arn MK TYPE REFERENCES v sin i REFERENCES ph (kni/sec)

37 120901-2 A2V F9II-III 120: £ 50 0,2

38 126269-70 A2V G1II-I11 80 £ 50 0.1

39 144208-9 5983 A2V F7II1 60 £ 50 0.1

40 156015 64 0 7 a Her B G5III F2V: £ 50 ind. ind, G511I:+F2: Bidelman (1958)

41 157978-9 6497 AlV GOI1 see note 90: 0 .2: G2lb Kuhi (1963)

42 159870 6560 G5III A7V £ 50 60 1.2

43 169689-90 6902 G8III AOV < 50 60 0.9 G8IV Kuhi (1963)

44 169985-6 6918 59 Ser A A0V: c o m see note 90; 0.3 GOIII+A6V Bahng (1958) 177 Boyarchuk & Kopylov (1964) G8III Kuhi (1963) AOVn+GO M eisel (1968) 300 ;80: Meisel (1968)

45 175492-3 7133 113 Her G5III A2V £ 50 60 1.1 A5+G5III S lettebak < 25 S lettebak (1955) (1955) G4III+A6V Bahng (1958) Boyarchuk & Kopylov (1964) G5III Kuhi (1963)

46 175580 G5III A0: V £ 50 ^ 50 1.3 KOIII+A Bidelman (1957)

47 183912 7417 P Cyg A K31I B9.5V < 50 50; 1.0 K3III+B: Stebbins 8 < 25 S lettebak Kron (1956) (1963) K0II+B9V Bahng (1958) 48 184398-9 7428 K 2 II-III AOV £ 50 < 50 1.3 K2II-III+A Bidelman (1954) K2III Kuhi (1963) 49 184759-60 7441 9 Cyg A0:V F8III £ 50 £ 50 0,2 50 187076-7 7536 & Sge M2II AOV < 5 0 £ 50 1.0 M2lb*II+A Bidelman (1954) M21I;+B: Stebbins 6 Kron (1956) M2II+AOV Bahng (1958) K2lb Kuhi (1963) 51 187321-2 B9V C5II £ 50 < 50 0,2 G SII-III+Aj B Bidelman (1957)

52 187949 7571 AlV f4 III : 150 50: 1.6: APPENDIX A; Continued 61

Previously Previously mk T m : v siu i (km/sec) Determined Determined nami SMC. PRIMARY SEC, Am. MK TYPE REFERENCES v sin i NO. 111) HR -: PRIMARY ph REFERENCES (km/sec)

53 192577-8 7735 31 Cyg A K?ll B4V £ 50 80 1.0: K 2 iw ;3 v Bidelman (1954) < 25 S lettebak (1963) K21I Kuhi (1963) < 25 Boyarchuk & Kopylov (1964) 54 192909-10 7751 32 Cyg K3J1 B9:V 4 50 85 1,5 K 3IH 1+E Bidelman (1954) < 25 S lettebak (1955) K516-II+ S le tte b a k (1955) K31I Kuhi (1963) 55 193495-6 7776 f) Cap A AOV: G5II see note 50 0,5: G8II Kuhi (1963) 0 Boyarchuk 6 Kopylov (1964) A0V:+G5III+ M eisel (1968) 120; 50 Meisel (1968) 56 194359-60 GOIII A3V 85 90 0,0 57 196088-9 B9.5V F4I1I £ 50 £ 5 0 0,5 58 196093-4 7866 47 Cyg K2lb B3V £ 50 £ 50 1,4 K4lb+ _ Slettebak (1955) < 25 S lettebak (1955) K2lb+B Bidelman (1957) < 25 Boyarchuk & Kopylov (1964) K2lb+B5V Bahng (1958) K2lb Kuhi (1963) 59 196753-4 7895 KOII-III A3V £ 50 80 1.4 KOI 11 Kuhi (1963) 60 200031 G51II A5V £ 5 0 £ 5 0 1.6 G5III+A Bidelman (1957) 61 202447-8 8131 a Equ C2II-III A4V < 5 0 £ 5 0 1.0 A+G5III S le tte b a k (1955) < 25 S lettebak (1955) G0III+A5V Bahng (1958) 25 Boyarchuk & Kopylov (1964] GOII Kuhi (1963)

62 202710 KOI II A7V £ 5 0 75: 2.0 KOIII+F Bidelman (1957)

63 203338-9 8164 Mllbep B2pe £ 5 0 50 0.1 Mleplb+B Bidelman (1954) Mleplb+B2 Stephenson (I960) Mllbep Kuhi (1963) Mleplb+Bl-.V Simonson (1968)

64 205114-5 8242 C2II B9V £ 5 0 £ 5 0 1.0 G2Ib+A,B Bidelman (1957) G8III Kuhi (1963)

( 62 AITCDJX A: Con tinned

Previously Previously MK T v r t - v sin i (km/sec) Determined Determined kami PKIMAIiY SIX!. Pi;m\i;Y sec, Am MK TYPE v sin i ffin-,i:L::ci.s NO. Ill) m: : ph (Km/see)

65 202216 8329 A2 V GO: 111 100 £ 50 0,8 Am Valkcr (1906) Haller (1966) k m Cowley (1967)

66 208233 A2V G 5III < 5 0 £ 50 0.7 m i Kuhi (1963)

67 208816 8383 VV Cep M2laep B < 50 ind, ind, M2cpla Bidelman (1954) M2+Ia-Iab Keenan (195?) B2;pe+M2epla Simonson (1968)

68 213310-1 8572 5 Lac MOII B8V < 50 90: 1,6 MOIb-II+A Bidelman (1954) MOIab+B Stebbins 6 Kron (1956) K5lb+A0 Hynek 6 Stanger (1959) K5lb+B?V Bahng (1958) MOIb Kuhi (1963)

69 214558 8617 C2III A4v £ 50 i 50 1.3

70 215182 8650 rj Peg C8II FOV < 50 50; 2,5: G8II: +F Stebbins 6 Kron < 25 S lettebak (1956) (1955)

71 215318-9 G1II-III A2V 70 120 0.6

72 218634 8815 57 Peg M 4lll A2V £ 50 90: 1 . 8: M4III+A3V Hackos 6 Peery (1968)

73 219512 A7V F3III; 90 80: 0 .4 ;

74 223047 9003 f And G5lb AOV < 5 0 £ 50 3.0: G5lb Roman (1952) < 20: Boyarchuk 6 Kopylov (1964 G5lb Kuhi (1963)

ind, = indeterminate. 63

Notes

1. In cases such as this, with an F-type secondary and a giant G or

K-type primary, the composite nature of the spectrum is difficult to ascertain. The only notable pieces of evidence are: Slightly abnormal strengths of the higher Balmer lines, especially H 12, 13, and 14, a slight veiling of the lines redward of 4100 X, and a slight weakening of the K-line compared to that of the G-type primary

standard, due to the Mfilling-in" by the continuum of the primary.

There is also a hint of a core in the K-line.

2. Narrow K-line core visible against wider washed-out background.

The background is narrower, however, than in most composites, in­

dicating an earlier type for the late-type component. This is

consistent with the G band strength and the general metallic line

strength, especially between H|3 and Hy. SB(20) .

3. Characteristic composite spectrum with a sharp K-line superposed

on a wide and very diffuse background, veiled metallic lines longward

of 4000 X, and strong Balmer lines shortward of the K-line.

4. Characteristic composite spectrum. The late-type is very similar

to 12 Comae (see number 35), but the early-type is somewhat earlier.

Wilson (1950) reported this star to be single.

5. Late-type component is heavily veiled. No evidence of diffuse

K-line background.

6. K-line region is very hazy with a very faint core visible. The

late-type spectrum shows almost no veiling. 64

7. K-line region is hazy with no core visible. The metallic lines

are moderately veiled throughout the spectrum.

8. Composite K-line with a very sharp core. Late-type component is

fairly heavily veiled. SB(20)„

9. Characteristic composite spectrum, K-line core visible but hazy

background is very faint.

10. K-line core is difficult to discern against the diffuse background.

Wilson (1941) noted this star as single with two observations. SB,

P = 1515.6 days.

11. Spectrum similar to T Persei (number 10); early-type is a little

earlier. Veiling is less than in number 10 because Am is smaller.

Wilson (1941, 1954, 1955) reported separations of 0V07 with Am =0.5

and < 0V08, and, also, single. SB, P = 5350 days.

12. Late-type component is very difficult to classify. From Hy to the

plate limit (about 5000 X) it seems to resemble a K1 H I standard, but

shortward of Hy, especially around H6, the metallic line spectrum is

much weaker, resembling a late F or early G-type spectrum. Calcium

emission is present at the K-line. SB(31).

13. As mentioned in note 1, this combination is difficult to recognize.

The evidence for a composite nature is: veiled metallic lines, K-line

core superposed on a wider, washed-out background, and the Balmer lines

shortward of He are stronger and wider than would be the case for a G

star alone. The exact spectral type of the early-type component is 65 difficult to discern because of the strength of the overlying metallic line spectrum in the region between H€ and the Balmer limit. High velocity star.

14. Barely discernable core at the K-line superposed on a faint, hazy, and wide background. The Balmer lines are weak.

15. Similar to HD 20084 (see note 13). K-line core is strong and wide. SB(38).

16. Faint K-line core visible against a very faint, hazy background.

17. Hazy K-line region with barely discernable core. This is due to the large ^ of the system. SB, P = 1576.44 days.

18. Characteristic composite spectrum with a very faint, sharp K-line core visible against a hazy background, The K absorption is reportedly always bordered by narrow emission components which vary slowly in re­

lative intensity. On this plate V > R. SB, P = 10,470 days.

19. Fuzzy K-line region with no core present. Pure B-type spectrum shortward of 4000 X with K-type spectrum clearly visible longward of

that. Strong absorption lines are seen at \3906 (due to Si I), X.3900 and just redward of Hll, in the B-type spectrum. Algol type variable

5™0-51f6 in 972.176 days. Lee and Wright (1960) report Am = 1.9 while

Popper (1961) finds = 2.2. SB, P = 972.16 days.

20. Sharp, narrow K-line core superposed on wide hazy background.

Balmer lines are very strong.

21. Characteristic composite spectrum. 66

22. Late-type spectrum heavily veiled. Very faint K-line visible with a sharp emission component on its red edge. There is no hazy background visible'. ~ Calcium emission component also present at red edge of H. Possibly weak Balmer emission.

23. Characteristic composite spectrum,

24. Little veiling of late-type component.

25. In the photographic region the B-type spectrum is very dominant, almost completely masking the late-type component. Emission appears at HP, Hy, and possibly at H5, but the line widths are not what one would expect for a star rotating near break-up velocity. This could be explained by assuming that we are viewing the system at or very near pole-on. The barely visible metallic lines of the K star appear to be possibly double but Batten (1967) has rejected this star as a spectroscopic binary. Miss UnderhiLl (1954) has done a thorough high dispersion study of this star using plates sensitive into the visible r e g io n .

26. Ultraviolet region is slightly under-exposed making the type of

the early component uncertain.

27. Characteristic composite spectrum.

28. W Cep type binary with general magnetic field present. The

Jascheks (1963) have given details of the variations in the spectrum.

On this plate the Balmer lines are visible at least to H17 and are

rather strong and broad with sharp emission superposed just violet of 67 the center. The K-line region, shows a faint, broad absorption back­ ground with a sharp emission line just violet of the center; no absorption core is present. Also present in emission are strong lines at \4068 [S II], \4243, 44 [ Fe II], \4287 [Fe II], X4358,9 [Fe II], and

\4413,16 [ Fe II]. Present but weaker are X.4134 and X4190 [Fe II].

SB, P = 9752 days.

29. Late-type spectrum strong, SB(32).

30. Very similar to HD 34318-9 (see note 20). The K-line core is difficult to resolve against the background. Sirius group. SB(29).

31. Late-type component is heavily veiled. K-line is sharp with no background. SB(51).

32. Slettebak (1963) has classified this as an Am star. Contrary to this classification, the spectrum appears to show the character­ istic composite spectra with a faint K-line core, fuzzy background, characteristic Balmer decrement for HJ3, Hi, HS and He, strong higher

Balmer lines, and veiling of the metallic lines between Hfi and about

4100 X.

33. The sharp K-line core is near the violet edge of a diffuse

background. SB, P = 14.50 days.

34. Luminosity late-type component is very similar to that of

12 Com (number 35) but the veiling is not as pronounced. SB,

P = 71.70 days. 68

35. K-line core visible. General veiling of late-type component features throughout spectrum, especially evident from 4400 X - Hf3.

Coma Ber. cluster. Herbig and Turner (1953) determined Am^ 0*

SB, P = 396.49 days.

36. Fairly heavy veiling of late-type component.

37. a.4077 somewhat veiled but ^4215 strong.SB(41) .

38. Weak veiling of late-type spectrum.

39. K-line core visible. Some veiling of the later type spectrum especially in the 4400 X - Hp region. Petrie (1943) found Am = 1.17, the A-type component fainter than the later type . SB, P = 108.08 days .

40. As with others of this spectral type combination, G + F, the exact type of the early component is vory difficult to discern. SB,

P = 51.58 days.

41. Early-type component is the primary and has broad lines. These broad lines are probably due to the presence of a close companion, of the same spectral type, making this a triple spectroscopic system

(Batten 1967), and not to rotational broadening. The late-type secondary is heavily veiled throughout most of the photographic region except in a few places between Hfi and Hy. Calcium emission is present a t the K -lin e . SB, P - 1170 days and 3.76 days.

42. Hynek (1938) lists this star as a suspected composite. The com­

p o s ite n a tu re is confirm ed. The spectrum shows some v e i lin g of the

late type component. SB(32). 69

43. Characteristic composite K-line. SB(51),

44. The K-line is faint and hazy, the Balmer lines are strong, and

the G band is weak. The latter is probably due to heavy veiling of

the late-type spectrum. The presence of a very close companion of

the same spectral type as the early-type component probably contri­

butes to the breadth of the lines and is not the result of rotational

broadening (see number 41). Wilson (1952) finds a separation of

0706: with one observation. SB, P = 386.0 days and 1.85 days.

45. Characteristic composite spectrum. The K-line does not appear

double as Hynek (1938) noted. SB, P = 245.3 days.

46. Very wide and fairly strong K-line with a faint core present.

Much wider and stronger K-line than in 113 Her (number 45).

47. K-line is broad and shallow with a trace of a sharp core. The

higher Balmer lines are like those in a B-type spectrum and there is

a very weak X.3820 He line visible. Some veiling of K spectrum present

from Balmer limit to just redward of Hy. Wilson (1950) finds this

star to be single.

48. Possible calcium emission at H and K. SB, P = 108.57 days.

49. K-line is hazy; the core is barely distinguishable. The higher

Balmer lines, H9 and up, more closely match that of an A2 V standard.

The late-type spectrum is somewhat veiled. Wilson (1950, 1951, 1952)

reports a separation of 071 between the components but the position

angles vary widely. He quotes = 0.5. 70

50. K-line region is very hazy; no actual line is visible. The H line of Ca II is in emission. The Balmer lines are strong and in­ dicate an A-type early component but a faint He 3820 X line is visible.

Hynek (1942) notes the periodic disappearance of the sharp, early type

K-line and the appearance of H and K in emission. He also notes that on three plates H, but not K, is in emission. Wilson (1941) found it single with two observations. SB, P = 3725 days.

51. Faint K-line core superposed on hazy background. Late-type spectrum is fairly heavily veiled.

52. V 505 Sgr, an eclipsing binary. The lines in the F spectrum are all very heavily veiled. SB, P = 1.18 days.

53. No K-line visible. Possible calcium emission at H-line. The early-type component is not earlier than B3 and from strengths of

Balmer lines, He line at 3820 X, and the general enhancement in the

regions of 4026 X and 4471 X over that of a normal K2 II spectrum, the early-type component appears to be about B4 V. Algol type variable,

V 695 Cyg, P = 3803 days. Batten (1967) quotes = 1.7. Wilson

(1950, 1951, 1952, 1954, 1955) has measured separations of 0'.'04 and

0'.'06: w idely v a ry in g p o s itio n a n g le s and a lso judged the s ta r to be

single on various occasions. SB, P = 3784.3 days.

54. K-line region is very hazy. Ca II emission present at K.

Haziness extends from He to violet limit of plate near 3400 X. A lgol

type variable, P = 1148.0 days. Wright (1952) finds Am^ = 2.7.

Batten (1967) also lists =2.0. SB, P = 1140.8 days. 71

55. Composite K-line feature clearly visible. The fairly broad Balmer lines are probably due to the presence of a close companion of the same or nearly the same spectral type as the early type component (see number 41). Wilson (1941) has judged this star to be single or less than 0'.'08 in separation. SB, P = 1374.13 days and 8.68 days.

56. Characteristic composite spectrum. G-type spectrum is slightly v e i le d .

57. B-type spectrum heavily veiling F spectrum. The early type spectrum shows a narrow K-line core and the presence of the \3820 line.

58. K-line region is hazy; Ca II emission is present and represented by a sharp line at the K region and possibly at H as well. He line a t 3820 K is strong and the regions at \4026 and X.4471 in the K-type spectrum are enhanced. Wilson (1950, 1951, 1952) has judged the star in separate observations to be single and double with separations of

0706: and 0705:

59. Hazy K-line region with possibly a faint core. SB.

60. Strong K-line core superposed on wide, diffuse background.

61. K-line strong and sharp; no evidence of diffuse background but the Balmer decrement is clearly of a composite nature. The late-

type component is somewhat veiled. Wilson (1950) reports the star as probably single. SB, P = 97.56 days.

62. Strong K-line core. 63. This very interesting spectrum has sharp absorption K-line super­ posed on a very faint, hazy background. At the violet edge of the background is a sharp emission line. Emission is also present at H and there are numerous emission lines due to [Fe IIJ, most notably at

4244 X and 4287 X. The higher Balmer lines are broad and a strong He I

\3820 line is present. The star is a VV Cep type binary with a 10th magnitude, normal B3 V visual companion separated by 4.5", but the composite characteristics are not due to it. Wilson (1950) reported the primary to be single,

64. Characteristic composite K-line.

65. This star has recently been classified both as a composite and an Am star by separate investigators. The spectrum offers a good example of the difficulty of distinguishing between these two possi­ bilities, especially at low dispersion. Although the K-line does not appear composite, there is heavy veiling of the metallic lines. By comparison with artificial composite spectra taken for this study, this appears to be a composite with a very dominant early type com ponent.

66. K-line region shows a sharp core superposed over an extremely faint washed-out background. The late-type secondary is somewhat veiled and is much fainter than the primary.

67. This very complex spectrum has numerous emission features, many due to [Fe II], notably at 4287 X and 4243 X. The Balmer lines can be counted up to H24 or H25; show ultra sharp with violet emission edges. Hfi and Hy, however, show broad emission features with sharp central absorption cores. There are also broad emission features between H8 and H9, at about 3855 X and 3860 X. The K-line of Ca II is faint and hazy with a sharp absorption line on the violet side and this in turn has a sharp violet emission edge. Wilson (1950, 1951) reported separations of < 0'.’12, < 0V10 and 0V05 with widely differing position angles. SB, P * 7450 days.

68. Calcium emission is present at H and K. The K-line region is hazy and like H has emission edges. The K-line also has a very faint,

sharp absorption core. Most of the spectrum is rather veiled. Hynek and Stanger (1959), in their study of this star, find that "sharp H

and K lines, with emission edges, appear superposed on the already composite H and K lines." They state that their radial velocities

are erratic and may not come from the early type secondary, perhaps

arising from an expanding shell surrounding the primary. They

estimate = 3,0. Wilson (1950) reports as single. SB(32).

69. Strong, sharp K-line superposed on washed-out broader K-line.

The late-type component is somewhat veiled from about 4400 X redward

to the plate limit near 5000 X.

70. K-line region shows a broad core superposed on a broader, hazy

background. The core seems too broad for an A-type spectrum. Veiling

is present redward of 4500 X. SB, P = 818.0 days.

71. Characteristic composite spectrum with little if any veiling. 74

72. As reported recently by Hackos and Peery (1968 ), this is a composite star showing an M-type giant primary combined with strong

Balmer lines. The K-line region shows the characteristic composite nature but the fuzzy background is not as wide as would be expected

for an M-type. The later-type features are somewhat veiled compared with an M4 III standard. There is probably a faint calcium emission

lin e a t H.

73. K-line is sharp with no perceptible background but the Balmer

lines show the characteristic decrement and the metallic lines are

all strongly veiled.

74. Spectrum shows a fuzzy K-line region with a faint but sharp core.

Possible Ca II emission at K on the red edge of the absorption core.

Wilson (1950) reports as single. 7 5

APPENDIX B

TABLE OF SUSPECTED COMPOSITES FOUND NOT TO BE COMPOSITE

NO. HD HR NAME SPECTRAL TYPE

1 3883 178 Am 2 9811 A3 lab 3 12447 596 a Psc A Ap 4 17378 825 A5 la 5 27295 1339 53 Tau Bp

6 30453 1528 Am 7 33054 1664 14 Ori Am 8 37202 1910 C Tau B2 IVp 9 43244 2228 42 Aur A 7 V 10 51424-5 2599 Am

11 66068-9 Am 12 76369 3552 17 Hyd 8 Am 13 88849 4021 Am 14 107054-5 4680 A8 I I I 15 135774-5 Am

16 173654 7059 5 Aql A Am 17 179143-4 Am 18 186745-6 38 la 19 187982-3 7573 A1 lab 20 209790-1 8417 17 Cep A Am

21 217675-6 8762 o And B6p 22 221615 8940 71 Peg M4 I I I 76

Notes

1. Bertaud and Floquet (1967) list A3 and F2 as the K-line and metallic line types determined from their own observations.

2. Hardorp et al. (1959) objective prism survey at 580 X/mm.

3. Slettebak (1954) notes that the spectrum shows strong chromium and silicon.

4. Johnson and Morgan (1953) standard.

5. Aller and Bidelman (1964) give an extensive discussion of this manganese star.

6. Bertaud and Floquet (1967) give the K-line and metallic line types as A7 and F2, respectively. SB; P = 7.05 days.

7. The Jascheks (1960) give the K-line, Balmer lines, and metallic

line types as A2, F2 and about F2, respectively.

8. Well-known shell star. SB; P = 132.91 days.

9. A spectrogram taken for this study suggests that this star can be

assigned a single MK type, A7V, with v sin i = 250 km/sec.

10. Hynek (1938) lists this as a composite (F5 + A2). A plate taken

for this study shows this to be a previously unreported metallic line

star with K-line, Balmer line, and metallic line types of A2-A5-F6

and v sin i < 50 km/sec (Markowitz 1969).

11. The Cowley3 (1965) have reported this as an Am star. The author

finds the type to be A4-A9-F7, with v sin i = 50 km/sec. 77

12. Slettebak (1963) has previously reported this to be a metallic line star with A4-A5-FO types and v sin i = 50 km/sec. On a plate taken for this study, the author finds the types to be A6-A8-FO and v sin i = 50 km/sec.

13. Slettebak (1963) gives A9-A9-F5 for the K-line, Balmer lines, and metallic line types, respectively.

14. Appenzeller (1967) has classified this spectrum as A9.5 III.

Palmer et al. (1968) give it an A8 V classification with v sin i = 135 km/sec. On a plate taken for this study the spectrum seems clearly to be that of a giant, near A8, with v sin i = 125 km/sec.

15. Similar to number 10. The author finds the types to be A6-F0-

F9 with v sin i < 50 km/sec (Markowitz 1969).

16. Slettebak (1963) found the types to be A2-A2-FO.

17. Babcock (1958) called this an Am star. The spectrum closely resembles that o f t Ursae Major is. The types are A5-F0-F9, with v sin i < 50 km/sec.

18. Morgan, Code, and Whit ford (1955).

19. Morgan, Code, and Whitford (1955).

20. Slettebak (1955) noted an A2 K-line and F2 metallic line spectrum.

21. Slettebak and Howard (1955) noted that this is a shell star.

22. A plate taken for this study indicates that the star can be

assigned a single spectral type, in the photographic region at least.

Spectral type by P.Keenan (private communication). Wilson (1950) noted

it as single. 78 APPENDIX C

TABLE OF CLOSE VISUAL BINARIES SHOWING COMPOSITE SPECTRA

NO. HD HR NAME SPECTRAL TYPE

1 25007-8 1230 49 Cep gG8 + A7

2 33883-4 1701 I o i l l + A2

3 34029 1708 a Aur G8III: + Fill:

4 37269 1914 26 Aur G 5 I I I : + A3V

5 40369-70 2099 K2III + A5V

6 49618-9 2520 14 Lyn G4III + A2V

7 59067-8 2859 G8lb-II + B

8 63208-9 3021 82 Gem G2III + A4V

9 75098-9 GO + A2

10 85558 3909 X Sex A2

11 88021-2 G2III + A2V

12 166479-80 6803 F2 + AO

13 186203-4 7497 47 Aql dF3 + A3

14 186518-9 7508 K3II + A

15 195692-3 7849 G + A2

16 199306-7 F2 + A2

17 2"3973-4 8595 F2 + A5

18 218640-1 8817 89 Aqr sg G2 + A2 Notes

1. Am = 0.9, separation = 170.

2. Kuhi (1963). Am = 0.2, separation = 076.

3. Stebbins and Kron (1956); type attributed to Morgan. Interferometer

double, see Anderson (1920). SB; P = 104.02 days.

4. Bidelman (1958); Kuhi (1963) gives GOIII for the late-type com­

ponent. Am = 0.4, separation = 073.

5. Bahng (1958); Kuhi (1963) gives KOIII for the late-type component.

Am - 1.8, separation = 075.

6. Bahng (1958). Am = 1.2, separation = 079.

7. Bidelman (1958). Am = 2.0, separation = 170.

8. Bahng (1958). Am = 0.1, separation = 072.

9. Aan = 2.5, separation = 078.

10. Aan = 0.4, separation = 078.

11. A pi ate taken for this study shows the above classification,

Am = 0.1.

12. Am = 1*1* separation = 175.

13. Am = 1.2, separation = 077.

14. Kuhi (1963). Am = 3.5, separation = 076.

15. Am = 1.8, separation = 076.

16. = 0,3, separation ^ 072.

17. = 0.5, separation = 079.

18. Am * 0.9, separation = 076. 80

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