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UNITED STATES DEPARTMENT OF COMMERCE • C. R. Smith, Secretary NATIONAL BUREAU OF STANDARDS • A. V. Astin, Director

Wolf-Rayet Stars

Proceedings of a Symposium held at

The Joint Institute for Laboratory Astrophysics

University of Colorado

Boulder, Colorado, June 10-14, 1968

Edited by

Katharine B. Gebbie and Richard N. Thomas

JILA

Institute for Basic Standards

National Bureau of Standards

Boulder, Colorado 80302

Held under the joint sponsorship of The American Astronomical Society, Harvard College Observatory The Joint Institute for Laboratory Astrophysics Smithsonian Astrophysical Observatory

Partially Supported by the National Science Foundation

\) National Bureau of Standards , Special Publication 307 ,

» » »- Issued December 1968

For sale by the Superintendent of Documents, U.S. Government printing Office Washington, D.C. 20402 - Price $3
National Bureau ef Standards JAN 1 4 1369

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ABSTRACT

A symposium on Wolf-Rayet stars was held at

the Joint Institute for Laboratory Astrophysics on

the campus of the University of Colorado, Boulder,

Colorado, 10-14 June 1968. The Wolf-Rayet stars

represent the most extreme example studied of an interaction between aerodynamic motions and a radiation field to produce a high temperature, largescale plasma in a steady but non-equilibrium state.

As such these stars provide a perfect example of

the kind of gaseous ensemble that JILA was created

to study. In order to understand them, we require

a knowledge of gases with temperatures between 10** and 10 °K and differential velocities between 0 and

10^ km/sec. In particular we need information on radiative and collisional atomic cross sections for

a wide range of ions, on collective interactions of

ions and photons, on methods of diagnostic spectroscopy, and on velocity fields generated by convective, gravitational, nuclear, radiative, rotational,

thermal, and other instabilities. The material of

the symposium was divided into four broad topics:

the distribution, physical properties and evolutionary status of Wolf-Rayet stars; the detailed features

of their spectra; the interpretations of these features

and the models on which they are based; and finally

a survey of the material and ideas arising out of the

symposium itself. This volume contains the introductory summaries of each of these broad topics, together with an edited version of the discussions

which followed.

Key words: atomic cross sections, atmospheric

aerodynamics, diagnostic spectroscopy, non-equilibrium gases, stellar instability, Wolf-Rayet stars.

Library of Congress Catalogue Card No. : 6 8-62108
CONTENTS

Preface v Acknowledgments vii

List of Participants ix

Opening Remarks Richard N. Thomas 1 Prologue The Wolf-Rayet Stars: Informal

Perspective C. S. Beats 7

Part A. THE FEATURES OF THE SYSTEM OF WOLF-RAYET

STARS 21

Chairman: Cecilia Payne-Gaposchkin

Introductory Speaker: Lindsey F. Smith 23

DISCUSSION 65

Part B. A SURVEY OF SPECTROSCOPIC FEATURES OF WOLFRAYET STARS 101

Chairman: Jorge Sahade

Introductory Speaker: Leonard V, Kuhi 103

DISCUSSION 144

Part C. SPECTROSCOPIC DIAGNOSTICS, INTREPRETATION

,

AND ATMOSPHERIC MODELS 181

Chairman: Lawrence H. Alter

Introductory Speaker: Anne B. Underhitt 183

DISCUSSION 203

Part D. SUMMARY OF PROBLEMS, IDEAS, AND CONCLUSIONS

ON THE PHYSICAL STRUCTURE OF THE WOLF-RAYET

STARS 2 37

Introductory Speaker: Richard N. Thomas 239

DISCUSSION 251

iii
The Spectriom of V444 Cygni ; Freddo, F. A.; Ap. J. , £1, 515, 1936.

There was a remarkable pygmi

,

Who arrived from V444 Cygni.

He said, as he grope.

Through his old spectroscope,

"It's them emission lines that intrygmi !

from "Aerodynamic Phenomena in Stellar

Atmospheres - a Bibliography",

Technical Note No. 30 (1959)

U.S. Department of Commerce

National Bureau of Standards.

iv
PREFACE

The symposium on Wolf-Rayet stars, sponsored

jointly by the American Astronomical Society, the

Harvard College Observatory, the Joint Institute for

Laboratory Astrophysics (JILA) , and the Smithsonian

Astrophysicai Observatory was held at JILA, 10 to

14 June 196 8. It was one of a series of symposia

that JILA has undertaken to sponsor in collaboration

with other institutions on topics of current interest

in overlapping areas of aerodynamics, astrophysics,

atomic physics, chemical physics, and the physics

of high temperature gases.

The Organizing Committee consisted of C. PayneGaposchkin, K. B. Gebbie , L. Goldberg, and R. N.

Thomas. There were thirty-eight invited participants,

and a number of graduate students and scientists in

related fields attended as auditors. The symposium was modeled after the Cosmical Gas Dynamics

Symposia, which are sponsored jointly by the International Astronomical Union and the International

Union of Theoretical and Applied Mechanics. Each of

the four formal sessions consisted of a survey paper

followed by general discussion during which participants were encouraged to present their views and

results where relevant. There were no contributed

papers

The National Science Foundation made a generous

grant toward the publication of these proceedings

and other expenses of the symposium. The balance of

the costs were covered by JILA (the University of

Colorado and the National Bureau of Standards)

In all the work connected with the symposium, we have been wholly dependent on Robert N. Alvis, Executive Officer of JILA, Mrs, Robert J. Low, whose ex- perience in running symposia is incomparable, and

Judy Schlepp, our secretary for astrophysics. Assisting them were Anne Cannon, Cordelle Yoder, and

Mary LaCasse. All technical facilities were coor- dinated by Stuart Jordan, who was on loan to us from

the NASA Goddard Space Flight Center. He was assisted

by three JILA graduate students, Stephen Hill, David

Van Blerkom, and Kenneth Ziebarth. We are also indebted to the JILA administrative and technical ser- vices and, in particular, to William Kellet, Victor

Holliger, and Floyd Howerton. The cooperation of

Mr. J. K. Emery of the University of Colorado PubV
lications Services is gratefully acknowledged.

Alice Levine has overseen from start to finish

the editing and preparation of the camera copy for

these proceedings. To her professionalism and good

sense of fun, we owe not only the appearance and

timing of this volume but also our own enjoyment in

editing it. Although each participant received and

corrected a draft of his remarks, the responsibility

for the final editing is ours. For instant responses

to demands for additional figures and references, we

are grateful to many of the participants and also to

a number of their colleagues who did not attend the

symposium. We were fortunate in being guided to the

quick and inexpensive form of these proceedings by

the Publications Office of the National Bureau of

Standards and, in particular, by W. R. Tilley, J. E.

Carpenter, and Betty L. Oberholtzer. The typescript

was prepared by Paulina Thure.

1 September 196 8 Katharine B. Gebbie

Richard N. Thomas
ACKNOWLEDGMENTS

We wish to thank the publishers of the following publications and the authors of a number of papers for permission to use figures and tables.

Annals of the Astronomical Observatory of Harvard

College

Annual Review of Astronomy and Astrophysics (Annual

Reviews, Inc.)

Astrophysical Journal (The University of Chicago

Press)

Bulletin of the Astronomical Institutes of the

Netherlands (North-Holland Publishing Company)

'Canadian Journal of Research (Canada Department of

Mines and Resources)

I.A.U. Symposiiim No. 34 Planetary Nebulae (D.

Reidel Publishing Company, Dordrecht, Holland)

Monthly Notices of the Royal Astronomical Society

(Blackwell Scientific Publications, Oxford)

Publications of the Dominion Astrophysical Observatory (Department of the Interior Canada)

Publications of the Lick Observatory (University of

California Press)

M. S. Burland

H. D. Curtis

W. A. Hiltner

C. R. Lynds

vii

PARTICIPANTS

L. H. ALLER, Department of Astronomy, University of

California, Los Angeles, California

C. S. BEALS, Manotick, Ontario, Canada

M. S. BESSELL, Yerkes Observatory, Williams Bay,

Wisconsin

K. BRACHER, Department of Astronomy, Whitman College,

Walla Walla, Washington

J. CASTOR, JILA, University of Colorado, Boulder,

Colorado

A. N. COX, J-15, Los Alamos Scientific Laboratory,

P. O. Box 1663, Los Alamos, New Mexico

J. COX, JILA, University of Colorado, Boulder,

Colorado

W. DEINZER, JILA, University of Colorado, Boulder,

Colorado

S. I. GAPOSCHKIN, Harvard College Observatory, 60

Garden Street, Cambridge, Massachusetts

R. H. GARSTANG, JILA, University of Colorado, Boulder,

Colorado

K. B. GEBBIE, JILA, University of Colorado, Boulder,

Colorado

R. HANBURY BROWN, Department of Astronomy, The

University of Sydney, Sydney, New South Wales,

Australia

C. HANSEN, JILA, University of Colorado, Boulder,

Colorado

R. M. HJELLMING, National Radio Astronomy Observatory,

Edgemont Dairy Road, Charlottesville, Virginia

D. G. HUMMER, JILA, University of Colorado, Boulder,

Colorado

H. M. JOHNSON, Department 52-10, Building 202,

Lockheed Missiles and Space Company, 3251 Hanover

Street, Palo Alto, California

S. JORDAN, Solar Physics Laboratory Branch, Building

2, NASA-Goddard Space Flight Center, Greenbelt,

Maryland

R. KANDEL, NASA- Institute for Space Studies, 2880

Broadway, New York, New, York

L. V. KUHI , Department of Astronomy, University of

California, Berkeley, California

K. C. LEUNG, NASA-Institute for Space Studies, 2880

Broadway, New York, New York

D. N. LIMBER, Yerkes Observatory, Williams Bay,

Wisconsin

D. MIHALAS, JILA, University of Colorado, Boulder,

Colorado

ix
K. NARIAI, Laboratory for Theoretical Studies, NASAGoddard Space Flight Center, Greenbelt, Maryland

C. PAYNE-GAPOSCHKIN, Harvard College Observatory,

60 Garden Street, Cambridge, Massachusetts

C. PECKER-WIMEL GORDON, NASA-Institute for Space

Studies, 2 880 Broadway, New York, New York

N. G. ROMAN, NASA Headquarters, 400 Maryland Avenue,

S. W. , Washington, D. C.

J. SAHADE, Observatorio Astronomico , Universidad

Nacional de La Plata, La Plata, Argentina

R. E. SCHILD, Mt. Wilson and Palomar Observatories,

1201 E. California Blvd., Pasadena, California

T. SCHMIDT-KALER, 53 Bonn, Poppelsdorfer Allee 49,

Germany

L. F. SMITH, Department of Astronomy, University of

California, Los Angeles, California

T. P. STECHER, Code 613, NASA-Goddard Space Flight

Center, Greenbelt, Maryland

C. G. STEPHENSON, Warner and Swasey Observatory,

Taylor and Brunswick Roads, East Cleveland, Ohio

R. N. THOMAS, JILA, University of Colorado, Boulder,

Colorado

A. B. UNDERHILL, Sterrewacht "Sonnenborgh , " Der

Ruksuniversiteit , Servaas Bolwerk 13, Utrecht,

Holland

D. K. WEST, Code 613, Goddard Space Flight Center,

Greenbelt, Maryland

B. E. WESTERLUND, Steward Observatory, University

of Arizona, Tucson, Arizona

J. D. WRAY, Dearborn Observatory, Northwestern

University, Evanston, Illinois

M. WRUBEL, Astronomy Department, Indiana University,

Bloomington, Indiana

X
OPENING REMARKS

Richard N. Thomas

On behalf of its joint sponso^rs, the American

Astronomical Society, the Harvard College Observatory, the Joint Institute for Laboratory Astrophysics

(JILA) , and the Smithsonian Astrophysical Observatory, I welcome you to JILA to participate in this

working symposium on Wolf-Rayet stars. We bring together a number of people interested in these stars

to discuss what these objects are and the problems

we must solve in order to construct satisfactory

models of them. The last such gathering to deal ex- tensively with this subject was held in Paris at the

College de France in 1939. At that meeting WolfRayet stars were considered in the context of the more general discussion of phenomena in novae, P Cygni stars, WR stars, and white dwarfs. The theme

there was the possible linkage of these stars

through the phenomena attending an ejected envelope.

Only superficial attention was paid either to the cause of the ejection or to its aerodynamical implications for the atmospheric structure and spectra.

In the present symposium, we shall concentrate on the WR stars themselves, considering related objects

only where they contribute to our understanding of

WR stars. We shall, however, explore in detail any

suggestions for a self-consistent kinematic structure of the WR configuration.

In 1949 I suggested that we were entering a new era of stellar atmospheric models : under the

impetus of our then new understanding of the outer

solar atmosphere, we were possibly beginning to con- sider models with a significant supply of mechanical

energy, instead of models whose properties were fixed

wholly by radiative energy transport, hydrostatic

equilibrium, and local thermodynamic equilibrium.

The question was whether such a mechanical energy

supply would affect all, or only parts, of the atmosphere, and whether it would have a significant effect on the momentum balance. The suggestion was

that the Sun represented one extreme where the influence of the mechanical energy was only marginally

detectable in the visual disk spectrum, and where it

did not significantly perturb the momentum balance.

The WR stars possibly represented the opposite ex- treme where the mechanical energy dominated those

1
factors controlling the formation of the visual disk

line spectrum and greatly perturbed the momentum

balance. Between these extremes were those phenomena

in other stellar atmospheres that are anomalous in

terms of the classical atmospheric model. No general suggestions were offered on the source or structure of the mechanical energy supply. It was recognized that our biggest problem lay in the development of diagnostic methods by which we could infer

the actual situation in these atmospheres. Now,

twenty years later, with the impact of new observations and a considerable increase in the sophistication in our understanding of how to analyze a stellar atmosphere without imposing the classical

restrictions, it seems worthwhile to look again at

the WR stars.

When we narrow our discussion to WR stars, we must begin by establishing what we mean by the WR

class. Spectroscopically , the broad classification

rests on 4 features which describe what I shall call

hereafter a "pure" WR spectrum,

1. The spectrum consists almost wholly of

emission lines. When absorption lines occur, they occur as satellites at the violet edges of the emission lines.

2. The emission lines are very broad. Interpreted as Doppler-broadening , the widths correspond

to differential motions of some hundreds to thousands of km/sec and are not necessarily the same for

all ions.

3. The lines in any one star represent a wide

range of excitation and ionization. The excitation

level of the line spectrum is generally much higher

than that of the continuum as estimated from its

spectral energy distribution.

4. The spectrum falls into one of two groups.

Either it shows strong lines from carbon and oxygen,

or it shows strong lines from nitrogen. Both groups

show strona lines of helium dIus other, weaker lines.

I suggest we designate as "classical WR stars"

those whose spectra show unambiguously the four

features of the "pure" WR spectrum and no other

features

The broad category of "classical WR stars" is

thus divided into two groups , each of which is further

divided into subclasses on the basis of relative line

intensities. We then have the choice between two alternative physical pictures:

a. Is a star that exhibits the "pure" WR spectrum a distinct kind of star that exhibits this spec2
triom as a consequence of its initial mass, chemical

composition, association with other astronomical objects, and subsequent evolution that carries it

through the WR stage at some point? If so, is each

subclass associated with a distinct kind of star in

the same sense, or do the several subclasses represent different stages in the evolution of one kind

of star? I suggest the model represented by this

first alternative be called a "WR object".

b. Does the "pure" WR spectrum simply imply a distinct kind of atmospheric condition, a distinct

state of the stellar atmosphere which can be attained

by different objects with different chemical compositions, from different causes, and along different

evolutionary paths? In this case, the different

subclasses would represent variations on this excitation state. I suggest that this model be called

the "WR phenomenon".

There are other objects whose spectra closely

resemble the pure WR spectra except that they either

lack some of the necessary features or show additional features. Still other objects have at some phase of their observed lifetime shown a spectrum

containing WR features. I suggest we call all such

objects "quasi-WR objects" and refer to their spectra as "quasi-WR spectra". I use the term "objects"

rather than stars because the observations do not

always refer unambiguously to a single star. Some

quasi-WR stars have been identified as binary, and

superposed on a "pure" WR spectrum are features

thought to come from the companion. Sometimes a quasi-WR spectrum appears to come only from part of

a star, such as from an ejected shell or nebula; it

is then difficult to- identify particular features

with given parts of the object. At this stage in our knowledge, we cannot say that if such-and-such

a structural feature of a given object were or were

not present, a "pure" Wolf-Rayet spectrum would result.

Because of the variety of quasi-WR objects, I lean toward my second alternative - that the pure WR

spectrum represents a phenomenon rather than an object. I would suggest that the spectral features

result directly from a supply of mechanical as well

as radiative energy to the atmosphere and that the

range from pure to quasi-WR spectra corresponds to

differences in the quality and quantity of this

mechanical energy supply. But at the moment this is

sheer conjecture: We are here to examine the actual

body of knowledge as a basis on which to test such

conjectures

3
We need two kinds of investigations. On the one hand, we need a taxonomic study of classical WR

stars, covering those properties not included in

spectral classification. Do they have other features

in common, and is there a relation between the spectral subclasses and any other physical parameter?

On the other hand, we require a detailed study of

the spectrum: What do we see, and what combination

of physical circumstances could produce it?

TAXONOMY

An empirical spectral classification scheme

does not fullfil its purpose until the variation in

spectral features from one class to another can be

correlated with a variation in physical parameters.

In the classical LTE interpretation of the HerzprungRussell diagram, these parameters were temperature,

density, and chemical composition, which were supposed to depend, in turn, on initial mass, initial

composition, and evolution. In the WR stars and re- lated phenomena, which are obviously non-equilibrium

configurations, it is not so easy to identify the

physical parameters. The taxonomic investigations

must be made without preconception as to which parameters are relevant.

In order to decide whether the WR spectral subclasses represent different stages in the evolution

of a single type of object or whether they represent

alternative stages in the evolution of distinct types

of objects, we must know as much as possible about

their differential physical characteristics. The same is true for quasi-WR objects. Indeed, if we could simply establish whether such things as mass,

luminosity, and the size of the differential velocity fields are constant across the broad WR category,

we would have made a start.

Questions of distribution and association with

other phenomena are obviously crucial for obtaining

absolute luminosities and masses of these objects

and may give some indirect clues on chemical com- position and evolution. Of equal importance in

settling the physical significance of the WR classification are the statistics on quasi-WR objects.

The problem of the possible binary character of the

classical WR objects - and of quasi-WR objects such

as novae - is also pressing. The search for an as- sociation with nebulosity, which may bear both on

4
the problem of mass ejection and on a possible asso- ciation of spectral characteristics with planetary

nebulae, is critical. Finally, any kind of observation or statistic that can give information on ve- locity fields should be pushed to its limit.

THE SPECTRUM

In some senses it is easier to comment on the

interpretation of the spectrum than to speculate

on what causes the atmosphere to produce the spectrum. We think we know more about the state of the

atmosphere that produces the spectrum than we do

about what effects that state. The following are brief comments on the four broad spectral characteristics of the WR class.

1. Emission Lines

To produce a spectrum that shows only emission lines - or that shows only emission lines with

a few absorption edges - is not trivial. Let me re- mind you of the alternatives.

a. Intrinsic Emission Lines

(i) A Schuster-tyye mechanism : Recent

work by Gebbie and Thomas (1968) casts severe doubts

on the utility of this mechanism.

(ii) Fluorescent effects: To produce

all the emission lines by such effects is, in practice, impossible, especially in view of the low ex- citation of the continuum relative to the line spectrum.

(Hi) Chromosphere-corona mechanism

mechanical heating: This is a tempting possibility,

especially in view of the large line breadths which

suggest large differential velocity fields, hence

a potential supply of mechanical energy.

h. Geometrical Effects

Whether the WR spectrum can be produced wholly

by the geometry of an extended, expanding, turbulent

atmosphere remains to be shown. We must remember

that there is little evidence of extensive dilution

effects

2, Broad Lines

The most pressing question is: Can a simple.

5
expanding, optically thin atmosphere give the observed line profiles, or do we require a system of

random motions coupled with optical-depth and radiative transfer effects?

3. Range In Excitation Conditions

This point is closely linked to (1) , (2) , and

(4) : What are the basic physical conditions that

can produce such a range in excitation and ionization? Is stratification necessary, or can uniform

but inhomogeneous layers produce the spectrum?

Clearly, an empirical model of the distribution of

excitation would be most welcome.

4. The Two Speotval Groups

The outstanding question is whether the existence of these two groups requires a difference in

chemical composition, or whether differential excitation alone can be the explanation. In order to

incite further questions, let me remind you that C.

Payne-Gaposchkin classified the first solar rocket

spectrum as WC6, and that C. Pecker-Wimel reproduced

this spectrum by using standard abundances in a rough model of the solar atmosphere. Let me also

remind you of the current efforts by Paczynski and

his colleagues to produce chemical differentiation

in terms of mass exchange in close binaries. Unfortunately, Paczynski, who will be with us at JILA

next year, was unable to come in time for this symposium.

We have organized the symposium in such a way

that our information on these points will first be

summarized and then systematically discussed. Starting with a taxonomic survey, we will go on to summa- rize the empirical spectral features, and finally

examine those conditions responsible for the production of the spectrum. The important aspect of the

symposium will be the discussion not only of the

summary papers, but also of the other discussions.

There is plenty of time, both in formal sessions and

in free periods, for a complete airing of all topics.

6
PROLOGUE

THE WOLF RAYET STARS: INFORMAL PERSPECTIVE

C, S. Beats

Let me begin my remarks by thanking Dick

Thomas for inviting me to take part in this symposium. The invitation is the more appreciated

since it was accepted with some trepidation due to

the considerable length of time that has elapsed

since my active participation in research on WolfRayet Stars. I am naturally sensitive to a comparison with Rip van Winkle who slept for 20 years,

waking up to experience difficulties in communicating with his friends and relatives who had outgrown

him. Perhaps a kinder assessment of my present re- lation to the Wolf-Rayet problem would be an analogy

with the semi-hibcr"^ or denners of which the

black bear is the most notable example. Unlike the

true hibernator who falls into a deep sleep in the

autumn and spends many months in a state of complete

unconsciousness, the bear sleeps lightly during the

winter, is easily disturbed and on occasion wakes

up, emerges from his den and investigates the state

of the outside world.

In a similar way, due to the force of circumstances I have spent considerable periods of time

in a state of mental somnolence induced by the performance of non-scientific tasks, but have wakened

up from time to time and used the nearest library

to see what was going on in the scientific world.

During these excursions I have become aware that

work on the WR stars has been going forward and that

the progress has been such that keeping in touch and

catching up is a major problem for anyone with limited time on his hands. I have therefore welcomed

this symposium as an aid in the catching-up process

and can only hope that my talk, necessarily con- cerned mainly with the past, will not only have some value as history but may also have some relevance

to the problems of the present day.

In talking about Wolf-Rayet stars I shall refer

to the work prior to 194 5 as Phase I and the subsequent observational work Including that now going

on as Phase II. My talk will be mainly about Phase

I with a few tentative excursions into Phase II. I will also make a few references to a Phase III which

7
is looming on the horizon and about which I have

great curiosity but little knowledge.

Phase I. General Discussion of Wolf-Rayet Spectra

The first of these stars were discovered just over 100 years ago on objective prism plates. On

such plates the emission-line WR stars stand out

as conspicuously as a Marxist at a Republican con- vention or a giraffe in a herd of shorthorn cattle.

While the discoverers were not unaware of the scientific promise of these bizarre looking objects, they were completely outside the mainstream of astronomical research of that day, which was concerned with

the absorption-line stars. In addition to these con- siderations, the statistically insignificant numbers

of Wolf-Rayet stars resulted in a long period of

neglect during which these stars received the treatment normally accorded to a group of non-voters in an election year.

Figure I. Objective prism spectrum of WC star

in n Carinae region.

8
Thirty years after discovery this period of

neglect was broken by a brilliant series of visual

observations of Wolf-Rayet bands by Campbell using

a slit spectroscope at the Lick Observatory in 1894.

The next major break came in the early 1920 's when

J. S. Plaskett at Victoria brought the spectra of

Wolf-Rayet stars into the open, and it is fair to

say that ever since they have been objects of very

great interest to astronomical science.

Subsequent to and including the work of Plaskett,

slit spectra have demonstrated the following charac25^

20C

t\I

XX mm4650, 4666

l$Q \

v

\ X 4666

) sa> fMo MO 0 soe mo tm

SO

6c AO Ma ^4666

"BStT zm 335 Iw^ Jm 500 fOOO t500 1000 t

u

n

HD. XJ*« ~

04 im m$ /m soo o soo m9fi tSOO 2000 2i

% m ii p

MJ). X 5696

}:^1 L ""JI 4\ m

III

IV

Figure 2. Spectrum line profiles from several

WR stars.

9
teristics of Wolf-Rayet spectra.

(a) The emission lines are very wide, of the

order of 50 A. Both high and low pressure explanations have been offered for the width, but it now

appears that low pressure explanations based on ejection velocities hold the field.

(h) The range of central intensities is very

great (from 25 x the continuum to close to zero)

(o) The profiles of lines are usually rounded

but are sometimes flat- topped, and the flat- topped

lines offer interesting possibilities for interpretation of stratification, the scientific possibilities of which have not yet been fully exploited.

In this connection it should be noted that a flattopped line finds its normal interpretation in a restricted distribution of ejection velocities,

none of which approaches zero. It seems likely that

any given ion would exist only in one particular

region or stratum of the envelope, where the physical

conditions for its production are favorable. Thus

given sufficient resolution in the observations and

interpretai:ion , we should be able to relate the distribution of expansion velocities to the distribution of ionization in the envelope. By contrast,

the explanation of a rounded profile appears to re- quire a velocity distribution that ranges from zero

to the maximum velocity displacement from the line

center. So here we do not have such a clear cut

method for studying the relation between velocity

and excitation.

(d) Displaced violet absorption lines have

been observed on some but not all or even most lines.

Strong displaced lines are observed for He II and

C III lines, and some of the associated emission

lines have flat tops. No alternative to the ejection interpretation of violet absorption lines has

yet been found.

(e) The level of ionization indicated by emission lines is high including He II, C III, C IV,

N III, N IV, N V, 0 III, O IV, 0 V and O VI leading

to the conclusion that the stars are of high temperature. No discussion of quantitative temperatures is

given here since it will no doubt be dealt with by

later speakers.

(f) The existence of parallel sequences be- came obvious early in the game and its cause is

still under discussion. In my own observations a carbon-oxygen sequence appeared pretty well exclusive of nitrogen and vice versa. Now this is not

so certain, but the division is still very well

marked. Suggestions to explain it include abundance

10
imiiMiiMi

1 :| IP^l';Hin

Figure 3. Two WR spectra of comparable excitation but of parallel sequences: HD192163 - nitrogen

sequence (above) ; HD192103 - carbon-oxygen sequence

(below)

and excitation. I have an intuitive preference for

the former but am sure that these points will be discussed in more detail later in the symposium.

(g) Although the continuous spectra of WolfRayet stars sometimes appear weak or absent in underexposed spectra, they are in fact always present

and of course very important. It appears that even

yet our knowledge of the distribution of energy in

Wolf-Rayet continuous spectra is insufficiently known.

This is a good open field for young astronomers, and

it should be mentioned that observations of the farultraviolet taken outside the Earth's atmosphere are especially important. It is essential to have observations over the whole field from the far-UV to the

far-infrared in order to make suitable correlations

between the character of the Continuous spectra and

the observed level of excitation in the Wolf-Rayet

envelope

Phase II. Relation of Wolf-Rayet Stars to Other Objects.

(a) Absorption O's. These are very dissimilar

in appearance. The O's show incipient emission but

of much lesser intensity and smaller width. The absolute magnitude of absorption O's appears to be

somewhat greater. The big question is: Can we re- gard temperature as the most important distinction

between Wolf-Rayet and 0-type stars or are there

other more fundamental differences? Certainly later

speakers will be discussing this point. I believe

11
1 / W

1

J 1

39 AO ^1 41 A3 .ca A5 A6 ^7 50

Figure 4. Tracings of continuous spectra of

WR stars: BD 37°3821 (above) and BD 35°4013 (below)

that a careful comparison of the distribution of

energy in the continuous spectrum for the two types

of objects over the entire observable region is of

prime importance in solving these problems. Those

involved in space science who have the facilities

to observe outside the Earth's atmosphere can m.ake a decisive contribution here.

(bj Novae. Some photographs of the spectra of novae are very similar to those of Wolf-Rayet stars,

so much so that a second look is required to say

which is which. Actually there are easily observed

differences. The excitation of Wolf-Rayet stars is

normally greater, the strong nova lines being mostly

hydrogen. The displaced absorption borders of novae

lines are usually more clearcut and are often complex, suggestive of a series of absorbing shells.

Also the nova spectrum is normally variable whereas

the WR spectrum is normally not. Each spectrum in

general seems explicable only in terms of ejection

either of successive shells as in novae or of more

or less steady ejection as for the Wolf-Rayet stars.

Some old novae show Wolf-Rayet characteristics,

and it would be in order to study them in more detail, particularly the distribution of energy in

the continuous spectrum.

(a) The P Cygni Stars. These stars show emission lines with absorption on their violet edges.

12
Figure 5. Comparison of WR and Nova spectra:

Nova Aquilae (above) and HD192163 (below)

In general the lines are narrower and the excitation

lower than in Wolf-Rayet spectra. There exist some Wolf-Rayet stars, e.g., HD151932, which would cer- tainly be classified as P Cygni objects if there were no other class to put them in. In general (not

always) , the P Cygni stars are of rather high absolute magnitude. The a Cygni stars often show Ha in

emission with a violet absorption border, and when

this occurs we include them with the P Cygni stars.

I believe that it is important to consider the WR

and the P Cygni stars together insofar as their ejection characteristics are concerned.

The P Cygni star HD190073 has a very interesting calcium line which appears in emission with a double absorption border on the violet edge. In a nova this line would be considered as an indication

of successive shells, but such an explanation is

difficult for this star since this peculiar line has

remained relatively unaltered for a period of many

years. I regard it vital to solve the problem of

this line if we are to understand clearly the process

of the ejection of atoms from the surface of any star.

(d) Nuclei of Planetary Nebulae. The nuclei

of planetary nebulae are sometimes Wolf-Rayet stars

and, at least at first glance, it would appear nor- mal for a star ejecting atoms to give rise to a spherical nebula. This is especially true if, as now seems probable, not all atoms are ejected

with very high velocities. It would appear legitimate to speculate whether all planetary nuclei were

once Wolf-Rayet stars.

Planetary nebulae with 0-type nuclei are in

general of higher excitation than those with central

Wolf-Rayet stars. It is somewhat difficult to reconcile this with the generally accepted idea that the

Wolf-Rayet stars are of higher temperature. This

difficulty seems to remain even if we add the WolfRayet bands to the total nebular emission.

The spectra of Wolf-Rayet planetary nuclei are

13
2'5

1-0

H.D. 190603

Hoc /

HY

KM/SEC -200 0 -^200

Figure 6. P Cygni star HD190603.
Figure 7. Ca II H and K lines, HD190073.

not always well matched with those of their surrounding nebulae, as strong nitrogen lines are observed

in nebulae with carbon-oxygen nuclei.

It seems reasonable to regard the relation of

Wolf-Rayet stars to planetary nebulae as part of the

general Wolf-Rayet problem, each part of which re- quires a solution before we can feel confident of our knowledge of this class of stars as a whole.

Observations of the continuous spectra of planetary

nuclei are of obvious importance in this connection.

(e) The Sun, Recent scientific papers dealing

with the Earth's outer atmosphere treat the phenomenon

of the solar wind. Certainly its existence is well

established, and its influence on the Sun's environment is important. Recently there have been

references to what is called the stellar wind.

These indications suggest that the Sun and most if

not all stars eject atoms from their surfaces to a greater or less degree. If we refer to the Sun's

ejection of atoms as a wind or perhaps a light breeze,

then the corresponding phenomenon in the vicinity of

Figure 8, Profiles of Ca II H and K lines of

HD190073.

15
Figure 9. Specimens of planetary nebul

NGC 2610 and NGC 3242. H. D. Curtis.

16
a Wolf-Rayet star would be a violent hurricane.

It appears that under certain circumstances the

ejection from limited areas of the Sun's surface is

greatly increased in volume and velocity. Such temporary bursts of activity are associated with solar

flares (hot spots?) and the effects on the terrestrial atmosphere consist of violent magnetic storms

and the appearance of hydrogen lines in the spectriam

of the aurora. It seems quite possible that phen- omena of the general nature of solar flares multiplied

many times in intensity and occurring on the surface

of a star of high temperature could well give rise

to phenomena such as are observed in connection with

Wolf-Rayet stars. In this connection it may be pointed out that spectroheliograph observations of the

Sun reveal the presence of exceedingly complicated

velocity effects in which the motions are not always

in the same direction although at times very large

volumes of material appear to be leaving the Sun's

surface for outer space. Insofar as I am aware, the

origin of these sometimes very violent motions is by

no means fully understood. If this is the case for

a star observed under as favorable conditions as the

Sun, it is perhaps not surprising that the causes of

the ejection of atoms from Wolf-Rayet and P Cygni

stars are still obscure.

(f) Binary Stars, A high proportion of WolfRayet stars are binaries, and it has even been suggested that they are all binaries though this is open

to question. The most striking of these binaries

studied during the period designated as Phase I is

the star HD193576, and it will probably be a long

time before its possibilities as a source of scientific information will be exhausted. Previous to

1945, observations of HD193576 had suggested (1)

evidence of asymmetrical bands due to tidal action

or other causes; (2) evidence that the radius of

the Wolf-Rayet envelope is greater than the distance

between the two stars; and (3) since primary eclipse

occurs when the Wolf-Rayet component is between the

0-type component and the observer, it appears that

the average surface brightness of the Wolf-Rayet

object is lower than that of the 0-type star. This

again raises the question of the relative temperatures of the two objects and the distribution of

energy over the whole observable range of the con- tinuous spectrum.

Figure 10 shows a profile of Ha at a time close

to secondary minimum when the 0 star is between the

Wolf-Rayet star and the observer. It is clearly a case of asymmetrical emission with a peculiar central

17
Id

k:

Figure 10, Asymmetric profile of Ha, HD193576.

absorption line. It is important to ascertain whether

the absorption line is a self-reversal of the WolfRayet envelope or an aspect of the absorption spectrum of the 0-type star. It is also possible that

hydrogen is an important contributor to the intensity of this line in addition to He II. It would

appear that a set of high dispersion observations

of this line covering the entire period of the binary orbit could lead to very valuable information

concerning this interesting binary star.

There are numerous other Wolf-Rayet binaries

which are now being actively investigated and some

of these will undoubtedly be dealt with by speakers

at this conference.

ig) Quasi-Stellar Objects. These objects have

been regarded by many as the most important astronomical discovery of the twentieth century. All kinds

of interesting observations have been made of them;

everybody is interested in them and, whenever possible, everybody wants to get in on the act. Certainly I cannot claim any special knowledge of them

and I have no desire to rush in where angels such

as William Fowler, Thomas Gold, Fred Hoyle and Martin

Schmidt fear to tread. Nevertheless, I feel bound

to mention the fact that C. R. Lynds has published

a spectrum of PHL 5200 which shows the presence of

emission bands wider than those of most Wolf-Rayet

stars. These bands have strong absorption borders

on their violet edges which are suggestive of P Cygni

18
.i

» « - • >

I

1

(£)Aro ^/\fO m

^ 'sT (J) O O

CM CM K> 5t ^53- in o

I > > ^ > >

Z CO CO O O

CO

CM = CO

o z

Figure 11, P Cygni spectrum of PHL 5200

C. R. Lynds.

stars. Nearly twenty years ago I was engaged in a study of P Cygni stars which was partly a compilation and partly a record of personal research. If

at that time the spectrum published by Lynds had

been available, the object PHL 5200 would certainly

have been listed among the P Cygni stars. While it

would be going a little far to list it among the

Wolf-Rayet objects, the great width of the emission

bands (of the order of 7800 km/sec) would certainly

make it of interest in any discussion of the origin

of ejection velocities.

No doubt Lynds and his associates will themselves be developing the consequences of this re- markable discovery. I only wish to suggest that the

presence in this object of Wolf-Rayet and P Cygni

characteristics may well influence the conclusions

on the nature of quasars, if they are indeed a homogeneous group of light sources. It might,

for example, give some inkling whether these objects

are small, bright galaxies or some sort of gigantic

star-like objects. Also when the nature of quasars

19
becomes better known, it might be that in the long run they could give some help in elucidating some

of the characteristics of the P Cygni and WolfRayet stars.

Notes on Phase III

Phase I above, to which my discussion has

mainly referred, and Phase II, which will be discussed by later speakers, have been concerned with

observable phenomena in the atmospheres or envelopes

or outer parts of Wolf-Rayet stars. Phase III,

mentioned earlier in this discussion, is here arbitrarily defined as referring to the fundamental

physical or structural characteristics which cause

a star to produce emission lines of the Wolf-Rayet

or P Cygni type. This is a problem which, it

would appear, can only be tackled from the theoretical angle aided by the reservoir of observational

data that has been accumulating for the past hundred years.

It is my impression that liaison between theorists and observational astronomers has improved a good deal since my early years. I seem to remember

quite a number of cases where reputably good theory

was combined with doubtful observations and vice

versa. In the first two or three decades of this

century the lack of modern computers often made it

difficult for the differences between eminent theoretical people to be resolved leading to impasse and

a waste of energy in controversy. We surely have reason to hope that in the future great advances in

theoretical studies, aided by the remarkable new computing devices, will make possible (to use a con- venient bit of political jargon) a more effective

consensus of theory than has ever been possible in

the past. In addition, the great increase in the

number and quality of observations with giant new telescopes will result in a combination that

should offer great possibilities for scientific

advancement in this field. I would like to close

by expressing the hope that this symposium will

give a needed and substantial impetus toward that

end.

20
PART A

THE FEATURES OF THE

SYSTEM OF WOLF-RAYET STARS

CHAIRMAN: Cecilia Payne -Ga-poschkin

Smithsonian As trophy sical Observatory

and Harvard College Observatory

INTRODUCTORY SPEAKER:

Department of

of California

Lindsey F. Smith

Astronomy ^ University

at Los Angeles

I. Introduction 2Z

II. Classification 24

III. Luminosities 27

IV. Distribution 30

V. Association with OB Associations

and with H II Regions 36

VI. Masses 42

VII. Interpretation of the Classification

System 46

VIII. Binary Stars 47

IX. Evolutionary Status 49

X. Source of Instability 50

XI. Do the WR Stars Form a Class, in the

Sense Defined in Section I? 59

XII. Summary 61

References 6 2 DISCUSSION 65

21

I . INTRODUCTION

My purposes in this summary are: (1) to draw

together the available information regarding the

luminosities, masses, ages and distributions of the

Wolf-Rayet (WR) stars, (2) to determine what correlations exist between these properties, and (3) to see what deductions may be made regarding the origin and evolution of these stars.

When we talk about a "class" of stars, we are implicitly assuming that all stars in that class

share certain common characteristics. The stars

need not be identical in all their properties, only

in those that are chosen to define the class. Observationally the characteristics most commonly

used to define a class are the spectral features.

We assume that spectral similarity results from

similarity in more basic properties of the stars,

and we attempt to define the spectral criteria such

that this is so. A further division of the stars

into subclasses represents an attempt to delineate

groups of stars all of whose properties are nearly

the same.

I submit that the necessary intrinsic characteristics that we require for a meaningful class

are twofold: (A) the stars are at closely related

stages of evolution, and (B) the principal mechanisms responsible for the defining spectral characteristics are the same for all stars in the class.

By "closely related stages of evolution" I mean

either (1) that the stars represent a characteristic phase, e.g., immediately post-main sequence,

in the evolution of a group of stars whose other

properties (mass, chemical composition) may cover

a wide range, or (2) that the different subclasses

of the WR class may be arranged in an evolutionary

sequence such that each star passes through a series

of successive stages in each of which it shows the

characteristics of one of the subclasses.

As a working hypothesis, I define the WR class

as those stars whose spectra show broad (greater

than 4 A) emission lines of highly ionized helium

and either nitrogen or carbon and oxygen; in particular they must have the He II A46 86 line and a broad band between X4600 and A4670 made up of either

N III and N V or C III and/or C IV lines. I exclude

novae, supernovae and the nuclei of planetary nebulae. The subclasses are differentiated according

to the dominant ions and the degree of ionization

23
evidenced by the spectra. (This choice is made

because we know that, for main-sequence stars, a similar criterion gives us groups of stars with

closely similar temperatures.)

I shall now proceed on the assumption that I have a class and subclasses in the sense defined

above. But I will pause at intervals to see if

there are any data which contradict this hypothesis.

Let me say right away that I strongly suspect that

the WR stars, as defined above, may not be such a class. In particular I suspect that single WR

stars, if they exist, may have different evolutionary histories than do WR stars in binaries and that

we may have to define the class more carefully if

it is to fulfill criteria (A) and (B) above.

II. CLASSIFICATION

The mode of classification currently applied

to WR spectra was adopted by the lAU Commission 29

in 1938 (Beals 1938) . The spectra are divided into

two sequences, WN and WC, according to the dominance

of lines of helium and nitrogen or lines of helium,

carbon and oxygen. The separation of the two se- quences is fairly complete, but there are a few

stars that display intermediate type spectra. The

spectral sequences are subdivided by Arabic numerals

denoting the relative degree of excitation evidenced

by the spectra. The definition of the subdivisions

has been the subject of some disagreement over the

past few years. At the moment two sets of modifications of the original criteria have been suggested,

one by Hiltner and Schild (19 66) and one by the

author (Smith 1968a) . Fortunately there is a fairly

clearcut relationship between the three systems;

this is given in Table 1. Which system is more

closely related to the parameters controlling the

spectral differences and is, therefore, the most

useful, is not clear at the present time.

In this paper I shall use the revised system

defined by Smith (1968a) . This system is effectively an extension of the original lAU system; it

gives a finer subdivision of the classes (e.g.,

spectra in the lAU class WC8 are now called WC9

,

allowing spectra between WC7 and WC8 to be called

WC8) and includes objects of higher excitation

than were known in 1938. The sequences run from

WN3 to WN8 and from WC5 to WC9 . Spectra denoted

by the lowest Arabic numerals have the highest

excitation.

24
TABLE 1

TRANSFORMATIONS BETWEEN VARIOUS

CLASSIFICATION SYSTEMS

Smith Hiltner-Schild Beals

WN3

WN4 (+0B)

WN4-A

WN5-A WN5

WN4.

WN5

5 (+0B) WN5.5-A

WN5-B

WN5

WN6

+ OB WN6-A

WN6-B WN6

WJnd WN7

+ UB

(+0B)

T*TK7 ^ C TV WNo . 5-A

WN7-A WN7

WN8 (+0B) WN8-A WN8

WC5

WC6

WC5

WC6 WC5

WC7

WC8

WC7

WC8

WC7

WC9 WC9 WC8

The presence of a binary may be detected in

several ways: (1) We may see a spectrum which is

most easily interpreted as a superposition of the

spectra of two stars. (2) We may observe periodic

variations of the radial velocity. (3) We may ob- serve light variations due to an eclipse. The presence of any of these characteristics does not guarantee that the object is a binary since cases are known in which these phenomena can occur without

the presence of a binary (e.g., pulsating variables

show variations of light and of radial velocity)

A careful evaluation can usually distinguish between

the possible interpretations; however, it is traditional that the title "binary" be qualified by an adjective indicating the variety of the evidence.

Binaries identified by observations (1) to (3) above

are called respectively, "spectrum binaries", "spectroscopic binaries", and "eclipsing binaries".

Clearly for any given system we may observe more

25
than one of these phenomena, and accordingly the

system may be called a "spectrum and spectroscopic

binary" or a "spectroscopic and eclipsing binary".

Among the WR stars we may distinguish two va- rieties of spectrum binaries. The first includes

those whose spectra contain absorption lines (other

than violet absorption edges to the emission lines)

as well as the characteristic WR emission spectrum.

Such objects are consistently found to be spectroscopic binaries as well, and hence we are sure of

the correctness of the interpretation of the spectrum in terms of superposition of spectra from two

stars. The second variety comprises stars whose

spectra show a much stronger continuum with respect

to the emission lines (or equivalently , weaker emission lines with respect to the continuum) than do

other stars with qualitatively similar emission

spectra. We suppose that in these cases the absorption lines in the spectriam of the companion are masked by the emission spectrum of the WR star.

Confirmation in the form of observed radial velocity

variations is available for a few of the stars of

this variety; however in any particular case, cer- tainty of the correctness of the interpretation is

lower for these stars than for those in which a definite absorption spectrum is observed. Spectrum

binaries are denoted WR + OB or, e.g., WN5 + 07 if

the subclasses of the components are known.

A spectroscopic binary may or may not be a spectrum binary, i.e. , one star may be so faint that

it does not contribute significantly to the spectrum. In the case of a WR star, this means we may

find that the velocity derived from an emission line

varies , but that there is no evidence from the spectral appearance for the presence of a companion.

This clearly opens the possibility that all WR stars

may be binaries , but that in many cases the com- panion has much lower luminosity and mass than the

WR star, so that we have noticed neither its con- tribution to the spectrum nor the change in the ra- dial velocity of the WR star.

Similar considerations apply to eclipsing binaries. However in this case there is little con- fusion and I need not discuss it further.

I have taken some care to describe the nature

of our knowledge about the presence of binary stars,

because it seems that one of the most important

questions facing us is whether all WR stars are binaries . The classification system proposed by Hiltner

and Schild (1966) divides the WN spectra into two

subsequences , denoted A and B according to the

26
strength and width of the emission lines, and as- signs Arabic numerals to denote relative degree of

excitation within each subsequence. Membership in

subsequence B requires broad lines. The subsequence

contains only WN5 and WN6 stars. [Although Hiltner

and Schild include one spectrum that they classify

WN7-B (HD62910) , it contains strong lines of both

carbon and nitrogen (H. J. Smith 1955, L. F. Smith

196 8a) and does not, therefore, truly belong.] Only

one star in subsequence B shows any indication of

binary nature. Subsequence A contains spectra with

narrow emission lines. Many of these stars are obvious spectrum binaries or known spectroscopic binaries. The implication is that all stars in subsequence A are binaries. This implies that the

stars in subclasses WN5-B and WN6-B are intrinsically different from all other subclasses, presumably in that they are single. It implies further

that all the single WN stars are in subclasses WN5

and WN6. Personally I find these conclusions hard

to believe for two reasons: (1) there is_ one known

spectroscopic binary in the B sequence, and (2)

there is a very smooth transition from WN5-B to

WN6-A and from WN6-B to WN6.5-A. From Table 1 you

will see that I have classified these as WN5 , WN5 + OB, and WN6 , WN6 + OB, respectively. It is, however, a very remarkable fact that in the WN sequence,

broad lines are found almost exclusively among the

apparently single stars in subclasses WN5 and WN6.

III. LUMINOSITIES

WR stars are very rare objects. In the Galaxy very few are near the Sun. Thus determinations

of their distances and the amount of interstellar

reddening are extremely difficult. For stars in

the Large Magellanic Cloud (LMC) , these problems

are minimal; the distance modulus of the Cloud is

known to within ± 0.2 mag, and interstellar absorption is small. Observations of apparent magnitudes of WR stars in the LMC were first made by

Cannon (19 24) ; more complete photographic photometry was carried out by Westerlund and Smith (1964)

,

and narrow-band photoelectric photometry was done by

Smith (1968) . With one exception, which I will discuss below, no significant spectroscopic differences

have been detected between WR stars in the LMC and

in the Galaxy (Smith 196 8b) . Thus we may expect

that the absolute magnitudes that we derive for WR

stars in the LMC also apply to the WR stars in the

Galaxy. The exception referred to is the remarkable

27
fact that stars in subclasses WC6 , 7, 8 and 9 are entirely missing from the Magellanic Clouds, and

stars in subclass WN6 are rare or absent. I do not

believe that this fact invalidates the assumption

that the LMC stars are similar to stars in corre- sponding subclasses in the Galaxy. However it does mean that absolute magnitudes of stars in the "absent" classes must be derived from observations of

stars in the Galaxy, and the values are thereby

much less certain. Results are given in Table 2.

When determining the magnitudes of WR stars, we need to give some consideration to the wavelength

range measurea. 'me emission lines contribute significantly to the total energy output in the visible

wavelengths; thus we need to specify which emission

lines have been included in any given measurement.

The values in Table 2 are derived from narrow-band

photometry (Smith 1968b). For WN stars, emission

lines are effectively avoided; the absolute magnitudes so derived are close to the absolute visual

magnitudes that the stars would have in the broadband Johnson-Morgan system if the emission lines were absent. For WC stars the v-filter includes a medium strong emission band at about X5140 due to

C III, C II and 0 V. The colors of stars in the

various subclasses of the WN sequence do not differ significantly, and a mean for all subclasses

is given; this has been converted from the narrow- band system into the UBV system; again it refers to

the continuum without emission lines. Individual

emission lines have a considerable effect upon the

narrow-band b-v measures for the WC stars; hence

TABLE 2 ABSOLUTE MAGNITUDES AND INTRINSIC

COLORS OF WR STARS

Class M

V

S.D. Class M

V

WN3 -4?5 0?1 WC5 -4 . 4 WN4 -3.9 0.3 WC6 -4.4

WN5 -4.3 0.1 WC7 -4.4

WN6 -5.8 WC8 -6.2

WN7 -6.8 1.0 WC9 -6.2

WN8 -6.2 0.4

W B-V = -0?08 ±0^06

28
such b-v measures have no well defined physical

meaning and are not reproduced here.

When the values in Table 2 have been determined from observations of stars in the LMC , a standard deviation can be derived; these numbers

represent the intrinsic standard deviations of the

magnitudes and colors within the subclasses.

For the subclasses that are not represented

in the LMC, we can derive the absolute magnitudes

for the WN6 and WC7 stars with reasonable confidence from observations of galactic stars (see

Smith 1968b) . The mean absolute magnitude of WC6

stars has been assumed equal to that of WC5 and WC7

stars, since the latter are equal. The only subclasses for which absolute magnitudes are really

uncertain are WC8 and WC9 . These values depend

only upon observations of Yi and y 2 Velorum. Graham

(1965) has determined the distance modulus of Yi

Velorum by H3 and UBV photometry. If we assume

that Y2 Velorum is at the same distance as yi Velorum, 4 60 pc, and that it is unreddened, we obtain

an absolute magnitude of -6.6 for Y2 Velorum. If we assume further that Y2 Velorum is a binary star

containing a WC8 star and a normal main-sequence 07

star (M^ = -5.2), we obtain an absolute magnitude

of -6.2 for the WC8 star. We have no information

regarding the absolute magnitudes of WC9 stars and

have simply assumed that they are equal to that of

the WC8 star in Y2 Velorum. = -6.2 is surprisingly bright compared to = -4.4 for the WC5 , 6 and 7 stars, although some of the luminosity may

be due to the contribution of emission lines. However, adoption of this luminosity does generate distances for the WC9 stars that are consistent with

the simplest interpretation of their asymmetric

angular distribution (see Section IV) . One further

caution regarding the values in Table 2: the absolute magnitude for WN5 stars rests upon two LMC

stars which have been classified by photometric

criteria. The classification of at least one of

these stars should be checked by spectroscopic observations . Using the values given in Table 2 (Smith 196 8c) , Wallerstein (1968) has plotted distances versus reddening as derived for galactic WR stars. Compared to similar plots for stars in Hiltner's (1956)

catalogue, the values in Table 2 give many more WR

stars with large distances and low reddening. Wallerstein points out that this may indicate that

some of the stars have lower luminosities than those

given in Table 2 and suggests that such stars may be

29
progenitors of planetary nebulae. Alternatively,

all of the scatter may be due to the fact that stars

at large distances can only be observed if the reddening is small. I would add to this the fact that,

plotted on a distance scale, scatter at large distances becomes very large because errors in the distance due to the natural range of luminosity within

a subclass are constant in log D not in D. Earlier attempts to determine the absolute magnitudes of galactic WR stars (Roman 1951, Andrillat

1955, Onderlicka 1958) were handicapped by uncer- tainties in the distances and reddenings of the

stars; they gave rather fainter values than those

in Table 2. A more recent attempt by Rublev (1963)

to determine the mean absolute magnitudes of galactic WR stars gives values consistent with those in

Table 2. Graham (1965) determined distance modulae

for young clusters and associations containing WR

stars. Resulting absolute magnitudes for WN7 stars

are within the range observed in the LMC, but the mean value is fainter. Other WR stars in clusters

observed -by Graham were used to determine the absolute magnitudes given in Table 2 for WR subclasses

that are not represented in the LMC.

Given the absolute magnitudes in Table 2, we notice a surprising fact; the brightest absolute

magnitudes are associated with those subclasses displaying the lowest level of excitation in their

spectra. This is proven among the WN stars and may

also be true among the WC stars. If there is any

correlation between the excitation temperature re- sponsible for the emission spectrum and the effective temperature of the star, then the above observation implies that the WN7 and WN8 stars are considerably larger in size than the stars in the

higher excitation classes. It is also true that if

there is a correlation between excitation temperature and effective temperature, the fainter stars

will have larger bolometric corrections, and the

bolometric magnitudes of the various subclasses

could approach the same value. Thus one of the

questions I think we should ask is: What relationship, if any, do we expect between the effective

temperature and the excitation temperature?

IV. DISTRIBUTION

In the Magellanic Clouds, we believe (Westerlund and Smith 1964) that we have found all of the

WR stars. There are 58 in the LMC and only 2 in

the SMC. As has been mentioned in Section III,

30
subclasses WC6 , 7, 8 and 9 are missing from the

LMC, and WN6 stars are rare or absent. The WR stars

in the SMC are both spectriim binaries, WN + OB and

WC5 + OB, and both have somewhat peculiar spectra.

In the Galaxy there are 124 known WR stars

(Smith 196 8a) . They are strongly concentrated to

the galactic plane (Roberts 19 62) and are frequently found in young clusters and associations (Roberts

195 8) and in binary systems in which the other com- ponent is a young O- or B-star. These facts are con- firmed by observations of the WR stars in the Magellanic Clouds (Westerlund and Smith 1964) and imply

that the WR stars belong to the extreme Population

I. Roberts (1962) also demonstrated that galactic

WR stars are strongly concentrated in directions

along which we observe spiral arms (which is con- sistent with their population assignment) and that

they are not found in the quadrant centered on the

anticenter. Reddish (1967a, 1967b) has noted that

they share the latter property with young clusters

and associations containing dust-imbedded stars and

with Bok globules.

Narrow-band photometry is available for most

galactic WR stars (Westerlund 1966, Smith 1968b) in

the photometric system in which the absolute magnitudes and intrinsic colors were determined (Table

2) . We can, therefore, determine their distances

and plot their distribution on the galactic plane

(Smith 1968c) . The result is shown in Figure 1.

The distance of the Sun from the galactic center is

taken to be 10 kpc (Arp 1965, Schmidt 1965). New

galactic coordinates are marked at the periphery,

and the different subclasses of WR stars are represented by different symbols as indicated. We see

the WR stars to very great distances. This results

from their high luminosity and from their ease of

detection on objective prism photographs. Recall,

however, that Wallerstein (Section III) suggests

that some of the large distances may not be correct.

The distribution of WR stars shown in Figure

1 shows elements of spiral structure as anticipated

from the work of Roberts. It also shows that the

various subclasses are not distributed in the same

way. In particular the WC9 stars show strong con- centration towards the galactic center, and the WN6

and WC7 stars also show a tendency to concentrate

to the inner regions of the Galaxy.

The presence of distribution differences is

most easily verified in terms of the angular distribution. In this way we are free from effects of

possible errors in the distance scale. I have

counted the numbers of WR stars in each subclass

31
Figure 1: The galactic distribution of WolfRayet stars. New galactic coordinates are given at

the periphery. The sun is marked +S ; the galactic

center is marked +C. The symbols used for the various subclasses are shown on the right.

that are within 60° of the direction of the galactic

center, compared them with the numbers of stars in

the same subclass that are more than 60° from

£ll = 0°, and applied a test for the significance

of deviations from a common distribution. The procedure was repeated for stars within 45° of = 0°.

The differences corresponding to those noted in Figure 1 are significant at the 5% confidence level.

The WC9 stars are all found within 45° of = 0°;

most of the WN6 and WC7 stars are found within 60°

of £H = 0°; most stars of the other subclasses are found more than 60° from 5,^^ = 0°. (The concentration of the WC7 stars to the inner regions only just

satisfies the significance test. It is imperative

that this should be re-examined when the completeness of the catalogue is improved in the fainter

magnitude intervals.) If the stars are distributed

with radial symmetry around the galactic center, it

follows that all WC9 stars lie within 7 kpc of the

galactic center, that most WN6 and WC7 stars lie

within 9 kpc of the galactic center, and that most

32
stars in other subclasses lie more than 9 kpc from

the galactic center. No distinction has been made

between single stars and binaries.

Thus we find that the population of WR stars

varies from the inner to outer regions of the Galaxy. It seems significant that the subclasses that

are concentrated towards the galactic center are among those that are rare or absent in the Magellanic Clouds. It is true that WC8 and WC6 stars are also absent from the Magellanic Clouds. There are only three WC8 stars known in the Galaxy; we do not,

therefore, have any reliable information regarding

their statistical distribution. The WC6 stars are certainly not concentrated towards the galactic

center. Thus the correlation is not complete. However it remains true that the population of WR stars

in the LMC is very like that in the Galaxy outside

9 kpc from the galactic center, but quite unlike the

WR population inside that radius. The outer parts

of our Galaxy (more than 12 kpc from the galactic

center) show a very sparse population of WR stars;

the SMC is more like this region than any other in

the Galaxy.

There are other objects in our own and in other

galaxies that show variations in properties from the

inner to outer regions. In particular, the mean

period of cepheid variables in M 31 decreases from

17 days near the center to 7 days in the outermost

spiral arm (Baade and Swope 1965) . From the work

of Kraft and Schmidt (1963) and Kraft (1963) , it

appears that there is a similar variation in our

own Galaxy. The mean period of cepheids in the LMC

is close to that of cepheids in the solar neighborhood (Shapley and Nail 1948) , while the mean period

of cepheids in the SMC is closest to that of cepheids in the outermost parts of our Galaxy (Kraft

1963) . Similarly, the ratio of neutral hydrogen to

total density in the Galaxy increases from less than

1% at the center to about 9% in the solar neighborhood and then falls off again in the outer parts.

The LMC, with a ratio of 9% (Bok 1966) , again compares well with the solar neighborhood, while the

ratio of 2 5% found in the Small Magellanic Cloud

(Bok 1966) is greater than that found in any part

of the Galaxy. (See Table 3.)

Thus it appears that the WR population shares

in the effects of some overall change of properties

from the center to the edge of the galactic disk

and from galaxy to galaxy.

I suggest that all of these phenomena are due

to different rates of star formation in the various

regions. It is reasonable to suppose that star

33
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34
formation has proceeded most rapidly in the central

regions of our Galaxy. This would be expected to

result in a depletion of the interstellar hydrogen,

as observed, and would presumably result in an in- crease in the heavy element abundance in the interstellar medium and, therefore, in the young stars.

I suggest that in the galactic disk, the rate of

star formation and the resulting heavy element abundance in the interstellar medium decreases monotonically from the center of the Galaxy to the outer

edges. Reports by Kraft (1963) that the cepheids

in the Per-Cas arm appear to be underabundant in

metals and by Williams (1966) that cepheids in the

Sagittarius arm are richer in metals than those in

the Cygnus arm are consistent with this suggestion.

Similarly Arp (1962) concludes that clusters originating in the outer regions of the Galaxy have lower

metal abundances than those in the inner regions.

However Arp (1965) suggests that the metal abundance in the galactic center is only as high as that

in the solar neighborhood, whereas the present data

suggest that the metal abundance in the galactic

center is significantly higher than is found in the

solar neighborhood.

It is known from the work of Hofmeister (1965a,

b) that the properties of the cepheid population are very sensitive to the abundance of helium and heavier elements. Thus, while the exact effects are not

yet known, differences in the chemical composition

of stars in the various regions are clearly capable

of causing the observed phenomena.

No direct evidence is yet available regarding

the chemical compositions of the WR stars. However

the present data strongly suggest that the initial

chemical composition plays a definitive role in

distinguishing the subclasses among the WC stars.

Since most subclasses of the WN stars share a com- mon distribution, we conclude that nearly all WN

stars, regardless of subclass, share a common initial chemical composition. The exceptions are the

stars in subclass WN6 , which are more strongly con- centrated towards the galactic center, implying

that they have a higher initial heavy element abundance than stars in the other WN subclasses.

The hypothesis that the observed distributions

are due to variations in chemical composition has

the attraction of being the simplest hypothesis

that is able to explain all the observed phenomena.

It is, in fact, the only hypothesis I have been

able to think of that will explain this. According

to this hypothesis, the similarity of the LMC to

the solar neighborhood and of the SMC to the outer

35
regions of the Galaxy results from a similarity in

the primordial density in the respective pairs of

regions and thereby in the past and present rates

of star formation.

V. ASSOCIATION WITH OB ASSOCIATIONS AND

WITH H II REGIONS

Determination of the association of WR stars

with other components of Population I has been

hindered by lack of distances and consistent spectral classifications for most of the WR stars.

Association of WR stars with clusters and

associations was first studied by Roberts (1958)

,

who concluded that about 20% of all galactic WR

stars are in clusters or associations and that a higher proportion of the WN stars are found in

clusters and associations than of the WC stars. He

suggested that WN stars evolve into WC stars.

Reddish (1967b) notes that WR stars are usually

found in very young clusters and associations in

which he finds circumstellar reddening of the most

luminous stars. He concludes that the WR stars are very young objects.

Shajn and Hase (1953) and Bok and Wade (1955)

have estimated that 50% of galactic WR stars are associated with nebulosity.

Observing WR stars in the LMC, Westerlund and

Smith (1964) found a correlation between the luminosity of the star and the youth of the association

in which it was located. WR stars in the 30 Doradus nebula are the most luminous, followed by WR

stars in other young associations, followed by WR

stars in the field. They conclude that the most

luminous WR stars are the youngest and the least

luminous are the oldest; they suggest that this

represents an evolutionary sequence.

Since we now know that luminosity is corre- lated with spectral subclass, it is easy to anticipate Roberts' (1958) suggestion: (1) the ages of

stars within a given subclass are fairly uniform,

and (2) the ages differ from one subclass to an- other. The following re-assessment of the situation and the conclusions are part of my doctoral

thesis and have not been published elsewhere, I therefore give them to you in some detail.

Knowing the distance of the WR stars, we can determine with greater certainty than before, which

stars are located within known OB associations.

Unfortunately the distances and sizes of the OB

associations are still only poorly known. I have

36
used the catalogue of associations given by Kopylow

(1958) . This catalogue was chosen in preference to

others because it makes a concerted effort to distinguish the associations that are strung out along

the same line of sight. Since most WR stars are located at considerable distances from the Sun,

this precaution is imperative. To make a statistical examination of the frequency with which WR stars occur in clusters and associations, I imposed a limitation of 16.1 mag on the apparent distance

modulus to which the WR stars were counted.

The location of known WR stars with respect

to H II regions was determined in the northern

hemisphere from the Palomar Sky Atlas and in the

southern hemisphere from the catalogue and atlas

of Rodgers, Campbell and Whitoak (1960). The distances of H II regions are rarely known. A WR star

has been assumed to be associated with an H II re- gion if it is centrally located with respect to the

nebula or with respect to some substructure of the

nebula. The same distance modulus limit was imposed for the WR stars associated with H II regions

as for the stars in associations. The nebulae fall

into two distinct categories. First, there are large nebulae which are usually associated with an entire OB cluster or association and whose ionization is clearly due to the combined effect of many

stars. Secondly, there are small nebulae which I have called "ring" nebulae. Here we observe the WR

star at the center of an arc or arcs of nebulosity.

It appears to be the sole exciting star for the nebula, and its central location implies that the

shape of the nebula is generated by some action of

the star.

Table 4 gives the numbers of galactic WR stars

within the distance modulus limits, and the frequency with which those stars fall in OB associations or in one of the two varieties of nebulosity.

Similar data are given for the WR stars in the Magellanic Clouds, except that no distance modulus

limit is necessary and ring nebulae are too small

to be detected in the LMC on the plates available

to me.According to the customary philosophy, stars

that are found in young associations or in regions

of nebulosity are themselves young. It is clear

from columns 3 and 4 of Table 4 that stars in some subclasses are found in young regions much more con- sistently than stars in other subclasses. Thus we may order the subclasses of the WR stars according

to the relative ages of the stars. This ordering

is shown in Table 5. The WN7 stars are nearly al37
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38
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.H CO in 00 in 00 + u WN WN WN WN WN WC WC DM WC WC OB 39
TABLE 5 RELATIVE AGES OF WR STARS

{Age increasing downwards)

ml

WN6

WN5

WN3

WN8

WN4

All binaries?

WC5

WC6,7,8,9

ways found in OB associations that are immersed in

nebulosity; the associations in the Carina nebula

and in the 30 Doradus nebula and the association

SCO OBI are prime examples. Thus the WN7 stars are clearly the youngest stars in the class WR. They

are followed by the WN8, WN6 and WC5 stars which

are frequently found in OB associations but not in

the extremely young associations in which we find

the WN7 stars. Next in age are the WN5 and WN4

stars that are only occasionally found in associations. Oldest of all are the WN3 and the WC6 , 7,

8 and 9 stars that are almost never found in asso- ciations . Thus among the WN stars those with highest

excitation spectra are the oldest, while among the

WC stars those with highest excitation spectra are

the youngest.

There is some evidence that the spectrum binaries are all very young, in some cases younger

than single WR stars of the same WR subclass. This

is surprising and important, if true. Unfortunately

the distance determinations for spectrum binaries

are a great deal more uncertain than for single

stars, and this result must be regarded with caution.

It should be noted that the "age" of the star

has the normal connotation of the time since the

onset of hydrogen burning. The length of the WR

phase may, of course, be shorter than the age.

The ring nebulae require separate consideration. Johnson and Hogg (1965) were the first to

realize the significance of the relationship between the WR stars and the ring nebulae; they suggested that the shape of the nebula is due to mass

loss from the WR stars, i.e., that the ejected matter sweeps up the interstellar material in the stellar vicinity into a thin spheroidal shell.

I have searched the Palomar Atlas for nebulae

with sharp filamentary rings. I find that every

40
such nebula not previously identified as a planetary nebula or a supernova remnant surrounds a WR

star. In the cases of planetary nebulae and super- novae we believe that the nebula is due to ejection

of matter from a central star. It may reasonably

be deduced that mass ejection is required for the

formation of such nebulae. The presence of nebulae

around WR stars thereby represents the most con- vincing evidence available that at least some WR

stars lose a significant amount of matter sometime

in their lives.

There are now seven known ring nebulae associated with WR stars. Four of these are shown in Fig- ure 2. They are found only around WN5 , WN6 and WN8

stars. In fact every single star of the subclasses

WN5 , WN6 and WN8 that is not in an association and

is sufficiently close that we would observe a nebula if it were present (i.e., within distance modulus 16.1 mag) is found to have such a nebula.

We would not expect to observe a ring nebula

around a WR star in an association. An OB association will ionize the surrounding gas so that such

a nebula, if present, will not be observed. I conlo/arc RCW 58 NGC 3199

NGC 2359 NGC 6888

Figure 2. Nebulae associated with Wolf-Rayet

stars

41
elude that all WN5 , 6 and 8 stars are probably ca- pable of forming ring nebulae, i.e., they are losing mass at a significant rate or have done so at

a previous stage of their evolution.

What can we say about the other subclasses?

Nothing definite. There are three conditions for

the formation of an observable ring nebula. (1)

The WR star must be losing mass, and the matter

ejected must carry a considerable amount of energy;

(2) The star must be located in a region where the

density of the interstellar medium is reasonably

high; (3) The star must not be located in an association. If one of these requirements is not satisfied, it is not possible to say whether the others

are

The three subclasses which do have ring nebulae are all in the three youngest groups in Table 5.

Nearly all nearby stars in equally young sxibclasses are located in associations. The situation with

regard to binaries is not clear. Stars in the

"older" subclasses may be sufficiently old that they

have moved out of the interstellar clouds from which

they were formed. Thus in each case requirement (2) or (3) may be unfulfilled, and we can therefore say

nothing definite about the mass loss from stars in

these subclasses.

If we know the total mass of a ring nebula

(e.g. , from radio observations) and if we assume

(1) uniform ejection of matter at a known rate, (2) a velocity of ejection deduced from the violet absorption edges on the emission lines (Underhill

1968) , and (3) the effects of radiation pressure

etc. to be negligible, then the requirement of con- servation of momentum allows us to determine the

age of the nebula. Assuming rates of mass loss between 10"^ and 10"^ solar masses per year, Johnson

and Hogg (1965) obtained ages between 10** and 10^

years for two nebulae. Assuming rates of mass loss

between 10"** and 10"^ solar masses per year. Smith

(196 7b) found ages between 10** and 10^ years for

five nebulae. If the proposed mechanism is correct,

it follows that the age of the nebula is equal to

the age of the WR phase of the central star.

VI. MASSES

We are not very certain of the mass of any

Wolf-Rayet star. The normal method of mass determination, via spectroscopic and eclipsing binaries,

breaks down for these stars because the velocities

42
derived from the emission lines in the spectrum of

Wolf-Rayet stars do not usually represent the ve- locity of the center of mass of the star. This

has been emphasized by Smith (1967a) who points out

that since different emission lines in the spectrum

of a WR star can give velocity curves with different amplitudes, eccentricities and phases, we have

no way of knowing which, if any, of the lines best

represents the motion of the center of mass of the

WR star. It is also pointed out that there appears

to be an orientation effect; the derived mass ratios are consistently smaller for systems viewed

edge-on than for systems viewed nearly pole-on.

The effect is demonstrated in another way by Figure 3, which gives a graph of the values of %r/Mob

against corresponding values of M^r sin^i derived

A 4686 for WN stars V A 4686 for WC stars O A 4650 for WC stars

M

04

WR

MOB

01

MwR Sin

3:

12

I© Figure Z, The relationship between derived

mass ratios of eclipsing binary systems containing

WR stars and M^r sin^i. The latter is dominated by

the factor sin^i and is thereby a simple measure of

the inclination of the orbital plane to the line of

sight. The values are taken from Table 6. The emission line used in deriving the plotted values is

indicated.

43
TABLE 6 DATA RELATING TO

Spectral

HD MR Name X 152270 65 WC7 + 05-8 4650

186943 94 WN4 + B 4686

4058

±y \jy ±o Q Q TaTM a RJ-HQ i^t = WLN*± . o+uy • Dj.a 4 D OD

193928 108 WN6 4686

197406 113 WN7 4686

4 ODUI

*i D *± JJ 4058

^ Q O ^ U Q ftOD Wr* ft + HR 4 ft<^

*t O D Ul

*± I ^

193576 106 V444 Cyg WN5 + 06 46861

Ae.f\c\. 4oUd>

4619J

114 CX Cep WN5 4686

211853 116 eclipsing? WN6 + BOI: 4686]

4603>

4058J

214419 118 CQ Cep WN7 4686

References

1. W. A. Hiltner 1944, Ap . J,, 99^, 273.

2. W. A. Hiltner 1945, Ap . J., lOl, 356.

3. W. A. Hiltner 1945, Ap , J., 102, 492.

4. W. A. Hiltner 1948, Ap , J., 108, 56.

5. W. A. Hiltner 1949, Ap . J,, 95.

6. W. A. Hiltner 1950, Ap . J,, 112, Ml,

7. R. M. Hjellming and W. A. Hiltner 1963,

Ap, J,, 137, 1080.

for WR stars in spectroscopic binaries (i is the

angle between the axis of rotation of the binary

and the line of sight). The abscissa, M^r sin^i,

is dominated by the factor sin^i; the correlation

implies that the derived value of the mass ratio

of a binary depends on the inclination of the or- bital plane to the line of sight. Since for any

given system, the different emission lines can give

different values for the velocity amplitude, the

derived mass ratio is also dependent upon which

emission line is used to represent the motion of

the WR star; the emission line used in the orbit

44
TABLE 6 (cont.)

MASSES OF WOLF-RAYET STARS

References

Spectro-- PhotoWR^ OB WR f (M) M^r/M® graphic metric

0.24 1.59 4.29 2,10,13

0.28 6 .

13.12 2,13

3 . 36 3 . 92

0.16 0.51 2.* 85 12,13

4.94 2 0.06 13

0.09

0.24

0.24 7.95 21.6 >8.0 3,13 7 0.31 5.23 10.0 >5 .

0.39 9.74 12.75 10.4 11 5,8,9

5.39 4 4 U . J D 16.01 >10.1 2,13 / 4.38 1 6

8. G. E. Kron and K,. C. Gordon 1943, Ap. J., 97,

311.

9. G. E. Kron and K. C. Gordon 1950, Ap, J., Ill,

454.

10. O. Struve 1944, Ap . J. . 100 , 384 . 11. O. C. Wilson 1940, Ap . J., 91, 379.

12. 0. C. Wilson 1949, Ap . J., 109, 76.

13. K. Bracher 1967, Thesis, Indiana University.

determinations has been indicated for each point

in Figure 3. A clearer correlation is obtained

when values of M^r/Mq^ for the various systems are derived from the same emission line.

An explanation of these phenomena, in particular the inclination effect, may give us considerable insight into the mechanisms operating in the

WR atmosphere. Until we have such insight, our determinations of the masses of these stars will

remain uncertain.

However we must have some sort of working

figure in mind. In Table 6 I have collected the

45
most recent information regarding the apparent mass

ratios and masses of WR stars in spectroscopic binaries. The sources of the data are given. In

most cases the numbers quoted are those derived by

Miss Bracher (1967) , who has made spectrographic

observations of many of the systems and has re- computed orbital elements for others.

We have three systems that are eclipsing binaries as well as spectroscopic binaries; these

contain a WN6 star, a WN5 star and a WC8 star. The mass estimates for the two WN stars are nearly equal:

10.4 and >_ 10.1 M®. For the WC8 star we have

>. 5 or >. 8 M0 (depending on which emission line is

used) . The > sign is still included because although we know the inclination is nearly 90° for

eclipsing systems , the light curves are not determined sufficiently well to define the value of the

inclination, and we know only that the masses are slightly greater than the values given. Thus if we guess that we obtain the correct answers when

we obsetve the system edge-on, we may assert that

the WN5 and 6 stars have masses of about 11 Mq and

that the WC8 stars have masses of about the same

or slightly less than those of the WN stars. We

should, however, bear in mind that we have found

considerable differences between the ages and distributions of the various subclasses. It would be

unwise to assume, on the basis of the above data,

that, for example, all WN stars have the same mass

or that all WN stars are more massive than all WC

stars

A ten-solar mass star on the hydrogen-burning

main sequence has an absolute visual magnitude of

about -2.5 (Allen 1963). The absolute magnitudes

given in Table 2 for WN5 , WN6 and WC8 stars are -4.3, -5.8 and -6.2, respectively. Thus these

stars appear to be 2 to 4 magnitudes over-liaminous

for their masses.

VII. INTERPRETATION OF THE

CLASSIFICATION SYSTEM

What now can we say about the classification

of the WR stars? What parameters seem to be important in the definition of the subclasses?

The distribution of the subclasses suggests

that chemical composition is an important parameter;

it suggests further that the WC9 stars have the

highest initial heavy element abundance, followed

by the WC7 and WN6 stars, followed by the remaining

subclasses which all have the same initial composition.

46
We have also found a difference in age among

the subclasses. The most likely explanation is

that the subclasses represent the same stage of

evolution in stars of different masses. This is a situation with which we are familiar; the youngest

subclasses originate from the most massive stars,

which evolve the most rapidly.

Sequential evolution through the subclasses

seems unlikely. It would involve the evolution of

WN7 stars, which we always find within young associations, into WN5 stars, which are sometimes found

"alone" and still within neutral hydrogen clouds.

This does not seem possible.

Let us see if, qualitatively, we can account

for the diversity of subclasses via the two parameters, initial mass and initial chemical composition. I suggest, strictly as a. working hypothesis,

that the WC sequence depends primarily upon initial

chemical composition , with the WC9 stars having the

highest initial abundance of heavy elements and

the WC5 and WC6 stars the lowest.

Most of the WN stars share a common distribution pattern in the Galaxy and, therefore, a com- mon initial chemical composition. I suggest that

the WN sequence depends mainly upon the initial masses and present ages of the stars, with the WN7

stars originating from the most massive and the

WN3's from the least massive stars.

We have one exception to this scheme in each

sequence. Amongst the WC stars, the WC5 and WC6

stars share a common distribution and hence a com- mon initial chemical composition. However, the WC5

stars appear to be significantly younger than the

WC6 stars, and this difference must be attributed to

a difference in initial mass. Similarly among the

WN stars , the WN6 stars appear to be demarcated by

initial chemical composition, as indicated by their

distribution. These exceptions indicate that if

this scheme is approximately correct, then the effects of initial mass and chemical composition are intertwined in a moderately complex manner.

VIII. BINARY STARS

The question we must now ask explicitly is:

Are all WR stars binaries, or are some single stars?

First let me present to you what seem (to me) to be

the most relevant facts.

1. Many WR spectra are clearly composite. Of

those that appear to be single, a few are spectro47
scopic binaries. Is it possible that all WR stars

are binaries but that the companions of some are of

low mass and low luminosity and have no conspicuous

effect upon either the spectrum or the velocity of

the WR star?

2 . The WR components of systems that are spectrum binaries appear to be very similar in spectral

characteristics and in luminosity to the WR stars

that appear to be single (with the exception of the

WN5 and WN6 stars discussed in Section II) . I will

comment further only on the data relating to the

luminosities of these stars.

a. Among the WC5 stars in the LMC (Smith

1968b) the so-called single stars are nearly 1 mag

fainter than the faintest binary star. The decrease

in contrast between the emission lines and the con- tinuum is qualitatively consistent with the hypothesis that the WR components all have the same luminosity but that the luminosities of the OB components differ.

b. The data noted in (a) clearly imply

that a given subclass of WR star may be found with

OB-stars of a wide range of luminosities.

c. No inconsistencies have arisen from

the assumption that the luminosity of a binary system is the sum of the luminosities of the so-called

single WR star of the appropriate subclass and an OB star of the appropriate subclass. (See Smith

1968b.)

1. Evide-noe in favor of all WR stars being binaries

a. It is likely that the WR stars in

binary systems have undergone severe mass exchange.

If we also have truly single WR stars, then we have

two groups of stars whose evolutions must have been

vastly different. The similarities of the spectra

and luminosities imply, but by no means prove, that

all the stars in a given subclass have the same evolutionary history.

b. Binary nature and mass exchange give

us a credible mechanism for creating a small number

of very peculiar stars.

2. Evidence against all WE stars being binary stars

a. There is some evidence, as outlined

in Section V, that the spectrum binaries are significantly younger than the apparently single WR

stars in the same subclass. Could this be a matter

48
of degree rather than kind? Is it possible that a WR star with a more massive companion evolves more

quickly or evolves differently?

b. The WR stars associated with ring

nebulae all appear to be single stars. (Kuhi (1968)

has suggested that HD50896, associated with RCW 11,

is a binary, but this is not proven.) Could this

be a statistical accident or another effect of degree?

I leave these as questions. But I note that

our conclusion on this point will affect our attitude towards the classification system and its interpretation.

IX. EVOLUTIONARY STATUS

From the data summarized here, I think we would still regard WR stars as Population I stars

and young objects. However, while stars in some subclasses, e.g., WN7, are extremely young, we may

justifiably wonder whether the stars in some other

subclasses, e.g., WC9 , may be near the age limit

for what we generally think of as Population I. We may, I think, be quite sure that WR stars

are not pre-main sequence objects, or at least

that they are not all pre-main sequence objects.

The fact that stars in some subclasses are fairly

old indicates this, but we may argue more precisely

as follows. Our mass estimates for WN5 and WN6

stars are about 10 solar masses, yet these stars

are not always found in extremely young regions.

By contrast, T Tauri stars, which are believed to

be pre-main sequence, have masses of the order of 1 solar mass and are only found in the very youngest

regions. Since a 10 star will complete its main

sequence contraction much faster than a 1 solar

mass star, a pre-main sequence star of 10 solar masses should be found only in regions at least as young as those in which we find T Tauri stars.

It is possible that WN7 stars, which we find

in very young regions, are in a pre-main sequence

phase; however, this contradicts the assumption

that the WR stars are a homogeneous group, and I would suggest that we do not introduce that complication unless it is necessary.

The extreme youth of the WN7 stars indicates

that the stars are probably no older than immediately post-main sequences. I think that we may safely

assume that the WR stars are not on the hydrogenburning main sequence; so we must conclude that the

49
WR stars represent a phase in the evolution of some or all stars that occurs immediately after the

stars leave the main sequence.

X. SOURCE OF INSTABILITY

The high velocities observed in the atmospheres of the WR stars, together with the evidence

of mass loss, lead us to suspect that the stars are on the verge of instability. What can be the cause

of this instability? Several suggestions have been

made over the last few years.

Sahade (1958) and Underhill (1966) have suggested that the WR stars are akin to the T Tauri

stars and that their instability arises from a similar source whatever that may be. This possibility

seems unlikely in the light of evidence presented

above that the stars are in a post-main sequence

phase.

Limber (1964) suggests that the emission shell

is generated by forced rotational instability which

produces tangential ejection. The instability is

triggered by contraction of a rapidly rotating star.

I will leave it to Limber to tell us whether this

model can account for many of the properties discussed above.

Recently Paczynski (1967) has suggested that

the WR stars are a product of mass exchange between the components of a close binary. He suggests that they are helium burning stars situated near the helium burning main sequence, and that

they may be on the verge of pulsational instability

due to the temperature sensitivity of the rate of

helium burning. He has calculated the evolution

of a 16 M© primary to the completion of mass ex- change. Kippenhahn and Weigert (1967) have calculated the evolution of a 9 primary to its departure from the helium-burning phase; Kippenhahn

(196 8) has now completed the evolutionary calculations for a 25 Mq primary. The detailed correspondence between the models and the observed properties of the WR stars has been worked out by Kippenhahn and the author. In what follows, I adopt the

convention used by Paczynski, that the star that is

initially the more massive is called the primary

throughout.

The evolutionary tracks in the log L, log Tg£f

diagram for the 9 M© star (Kippenhahn and Weigert

1967) and for the 25 M^j star (Kippenhahn 1968) are shown in Figure 4. The mass of the secondary has

50
Figure 4, The evolutionary tracks in the log L

- log Tgff diagram for primary stars in close binaries. In both cases the separation of the binaries

is such as to give mass exchange after depletion of

hydrogen in the core (Kippenhahn and Weigert 1967

Case B) . The upper track for a 25Mq initial primary

is from Kippenhahn (196 8) . The lower track for a 9Mq primary is from Kippenhahn and Weigert (1967)

virtually no effect upon the evolution of the primary.

Consider the track shown for the 25 Mq star.

Point A corresponds to the turn-off point from the

main sequence; after this the star moves rapidly

along the path shown. If the star were single, it

51
would follow the dashed path to the red giant region. Phases at which the star is moving from

left to right are phases in which the envelope is

expanding. If the star overflows its critical

equipotential before reaching point A, evolution

proceeds as described by Kippenhahn and Weigert

Case A (1967) , and the result is an Algol-type system in which the primary, which is by then the less

massive star, is in the subgiant region and fills

its critical lobe, while the secondary, which is

then the more massive star, is on the main sequence.

Agreement between theory and observation for these

systems is good and gives confidence that this process does operate frequently and in the manner indicated by the calculations.

If the separation of the binary is slightly

wider than that required to give mass exchange before point A, but is close enough that the primary

fills its lobe before reaching the equilibrium helium-burning (red giant) phase, we have Kippenhahn

and Weigert Case B. In the case shown in Figure 4,

the separation was chosen so that the primary filled

its critical equipotential lobe at point B. Any

matter that overflows this lobe is transferred to

the secondary. The primary continues to expand,

and rapid mass exchange ensues; the luminosity drops

as energy goes into expansion of the outer layers of

the star. Mass exchange continues until the primary

reaches a configuration in which its equilibrium

size (i.e., the size at which helium burning begins

in the core) is equal to that of the critical lobe.

This will take place just before the star is stripped

to its helium core. As it approaches this configuration, mass exchange slows, and the star brightens to

approximately its initial luminosity. At point D

helium burning begins, the star contracts away from

the critical surface, and mass exchange stops. The mass of the star is now a little greater than the mass of the helium core at the time of depletion of

hydrogen in the core. As shell burning increases the size of the

helium core, the star moves to the left, eventually

settling down at point E where most of the energy

is generated in the helium-burning core. The star

is near the pure helium main sequence; the presence

of a thin hydrogen-rich envelope holds it to the

right of that sequence. Notice that the final temperature, 5.5x10 °K, is the same for the 8 . 5 star

that results from evolution of a 25 primary as it

is for the 2 star that results from evolution of

a 9 M@ primary! The final temperature is very sen52
sitive to the amount of hydrogen -rich material that

remains after mass exchange, and this is difficult

to calculate; thus the equality shown here may be a coincidence, but it does suggest that the final temperature may be rather insensitive to the mass of

the star.

Boury and Ledoux (196 5) have shown that a pure

helium star is pulsationally unstable for masses greater than 7-8 M®. The instability results from

the temperature sensitivity of the rate of energy

production from helium burning. The product of this

class of close binary evolution will always have a thin shell of hydrogen-rich material which will damp

the instability; however, it is expected that for a sufficiently massive star, instability will be present. Van der Borght is presently analyzing the

stability of such models; results should be available in the near future. If such an instability is

present, it may result in shock waves that propagate

outwards through the star, producing an inverse temperature gradient and high excitation, as observed

in the solar chromosphere.

At point E the system is detached (i.e., neither

component fills its critical lobe), and the radius of

the 8.5 M© helium star is about 3.5 solar radii. It

will remain here for about 5x10^ years. As the star

evolves away from the helium-burning configuration,

it expands and again fills its critical lobe, transferring still more of its mass to the secondary.

The star is then at point F, with an effective temperature of about 2.5xlO'*°K. Carbon burning will

start in the core, and there will be both a heliumburning shell and a hydrogen-burning shell. Thermal

pulses originating in the shell sources are expected;

these may cause mixing between the layers above and

below the shell. Pulsational instability due to temperature sensitivity of the carbon-burning process

will almost certainly be present. The star will not

be unstable to mass loss as in the evolutionary stage

B-C, so it may remain in this configuration, filling

its critical lobe, for a time comparable to the core- helium-burning lifetime. Lifetimes in each of the

evolutionary phases are given in Table 7. The only

phases long enough to be observed are E and F. The luminosity at point E is slightly greater

than that of a pure helium star of the same mass. The excess luminosity is contributed by the hydrogenburning shell. We may easily extrapolate the re- sults for the 8.5 and 2 M® stars to predict the final

luminosities for any final mass and thence for any

initial mass. This is shown in Figure 5. At the

53
TABLE 7 LIFETIMES OF A 25 PRIMARY STAR OF A

CLOSE BINARY SYSTEM

Lifetime Critical

Stage (years) Path lobe filled

H core burning 4,520,000

H shell burning before

mass exchange 190,000 A-B

Rapid mass exchange;

luminosity falling 300 B-C yes

Slow mass exchange;

luminosity rising 3,860 C-D yes

He core burning ; moving

towards He main sequence" 62,000 D-E

He core burning

Tgff decreasing 480,000 E-F

C core burning 480,000? yes(?)

bottom is plotted the final mass of the primary.

Since the final mass is essentially that of the helium core at the time of depletion of hydrogen in the

core, the initial mass has a simple relationship to

the final mass and is plotted at the top of the diagram. The circles represent pure helium stars after

Van der Borght and Meggitt (1963) . The diamonds represent the 2 Mg) and 8 . 5 M@ final products calculated by Kippenhahn and Weigert (196 7) and Kippenhahn

(1968) respectively. The curve is drawn through the

two evolved models, parallel to the curve defined by

the helium stars. Assuming the stars radiate like

blackbodies at 5.5xl0'*°K, we determine a bolometric

correction of 4.60 mag; absolute visual magnitude is

plotted at the right. Due to uncertainty in the

effective temperature and bolometric correction, the zero point of the absolute visual magnitude is un- certain by about one magnitude. The temperature

calculated is an upper limit; thus the absolute visual magnitudes given are fainter limits. Bolometric

magnitudes of hydrogen-burning main sequence stars

are fainter by about 3 mag. Due to the high tem54
OJ CO r— CO ID ^ s

to

to 'a (U 1 1 5-1 0) 1 (U iH to rH w to n3 rH 0 U (U fd 0 to CO -P rH U (U -P Q) 5-1 Ti u (U « w fd O to -H •H ^ > 0 H grH fd4-i0>iEHx; -P U -P Q) H-i- B (U US ^ ^ •H 5h -P 0

cn 0 0) C fd U rH

rH rd 0 •H -H ii: • 0 -P I) X

x; U .H cnE-i -P -P to u -P fd -p MH LT) (U a c: 0) 5-1 C -H <D Xi E MH . 0 C Ji: 0 (D 15 Pu-P Q) >i (U in

Eh c •H (U PQ • to 0 rH •H •H -p 0^ (U ^ cu to rH fd 5h 0) O 5-1 5-1 00 M .H fd ^ ;3 fd 0

CO ;3 0) >^ ci^ (U > fd u 0) to (U }^ fd c; -H 5h Q)

fd rH ^ 0 (U •H > -p 0 5h -p o — -p c: g M-l to o (U c fd to to to Q) 0) -p tn (U > -P tJ U -H ^ -P 0 T3 fd u c; to -p c (u •H (U M o Q) fd (U 5-1 -H 0 > U Q) o 0) rH c -P (U -P 5h (U tn e rH 1 > > 0 fd >iM-l o Cu-P cn fd 0 C 5-1 5-1 to >

-p fd X OJ 0) -H > H ^3 ^ •H CD •H OJ 5-1 fd S (U tr> u U < CP 0 -P 55
perature of the helium stars, their bolometric cor- rections are larger. For a temperature of 5.5xl0'*°K

the difference in visual absolute magnitude between

the helium stars and hydrogen-burning stars of the same masses is about 1 mag.

At phase F, the star is a little more luminous

than at E. However the temperature is lower and is

about equal to that of a main—sequence star of the same mass. Thus the visual magnitude of the star

is more than 3 mag brighter than that of a main-sequence star of the same mass. We can also predict the final mass and luminosity of the secondary star, which receives the mass shed by the primary and moves up the main se- quence, becoming the more massive component of the

pair. (The initial main sequence calculated by

Kippenhahn (196 8) has been used throughout these

calculations.) There is, at this point, one serious

uncertainty. Mass loss from the system as a whole

is likely, since the secondary may find itself with

an angular momentum problem (cf . , Paczynski and

Ziolkowski 1967). We have, at this time, no way of

estimating how much mass is likely to be lost from

the system. Figure 6 shows two possible cases. The

graphs show loci of points of constant AM, the difference in magnitude between the final stars, in the sense that if the helium star at phase E is brighter,

AM is negative. The initial mass of the secondary

is plotted on the vertical axis; the initial mass of

the primary, on the horizontal. The latter may be

replaced by the final mass of the primary, which is

given between the two graphs. The upper graph shows

the case of no mass loss to the system. For a bolometric correction of 4.60 mag, AM is always positive

in the mass range for which we have results. If,

however, we have overestimated the bolometric cor- rection by, say, 1 mag, then the locus given for

AM = 1 mag becomes the locus for AM = 0 mag, and

the two stars have equal visual magnitudes. Similarly, if we consider phase F, the bolometric cor- rection is 2 mag less than used to derive Figure 6,

and 2 mag should be subtracted from the values of

AM attached to the loci. The lower graph shows the

resulting loci if only half the mass shed by the

primary is captured by the secondary, and the rest

is lost to the system. In this case, AM is consistently smaller than in the upper graph, and the helium star is often the brighter of the two components. For higher masses of the primary, the luminosity differences tend to favor the helium star.

These, then, are the essential properties pre56
10 20

'^1<'"*^3*g" 40 50

No Mass Loss 15 4^ /

7 /

10

05

2 4 6 8 lb Mi(final)

12 14

507o Mass Loss

10

^/ /

/

^^AMy- 0

aMv -10 / _ - - —

30-

20 Mo

30-

20 M2i

10

10 20 30 Ml (initial) 40 50

Figure 6. The loci of AM (the difference in

absolute visual magnitude between the components of

the binary) when the primary is in phase E are shown

The upper graph is calculated on the assumption of

no mass loss from the system. The lower graph as- sumes that 50% of the mass shed by the primary is

captured by the secondary. The vertical axes give

the initial mass of the secondary. The horizontal

axes at top and bottom give the initial mass of the

primary, and the final mass of the primary is between the two graphs. When the main-sequence star

is the most luminous, AM is positive.

57
dieted for the models that result from mass exchange

based on the stated assumptions. Let me compare the

predictions with the features of the observed properties of WR stars summarized in this report. The

theory predicts:

1. Helium stars will occur in binary systems.

Unless there is severe mass loss from the system,

the helium star will be the less massive component.

The companion is expected to look like a normal

main-sequence star.

2. The helium star will have a fairly high

effective temperature and may have a heli\im rich

atmosphere. In the case of the 8.5 M® helium star,

the helium concentration in the atmosphere is 46.5%

by mass at stage E, and 65% by mass at stage F;

these are only slightly greater than normal. However, thermal pulses in stage F may cause mixing

of the outer layers and may increase considerably

the percentage of helium and heavier elements in

the atmosphere.

3. The stars are overluminous for their masses

by 3 mag in Mj^q^/ ^ "^^9 iri Mvis phase F,

and by >. 1 mag in M^ig phase E. From Figure 5 we see that the predicted absolute visual magnitude

for a 10 M<5 star at phase E is -3.5. This is about

1 mag fainter than most of the values given in Table

2. However, considering that the bolometric magnitude used is an upper limit, the agreement is satisfactory . 4. The initial and therefore the final mass

of the secondary component is a free parameter.

Thus a helium star of given properties can be found

with main-sequence stars having a wide range of masses and luminosities, as observed. The helium

star can be either the less luminous or the more

luminous star of the binary. This allows the possibility that many binaries may go undetected; i.e.,

the absorption- line spectrum of the main- sequence

star may be completely hidden by the stronger emission spectrum of the WR star.

5. The stars are in an immediately post-main

sequence phase of evolution.

6. Pulsational instability may exist; if so,

it provides a supply of mechanical energy to excite

an emission-line spectrum.

7. The final luminosity of the helium star

depends only on its mass, which depends only on

the initial mass of the star. The age of the star

is essentially its main-sequence age, which also

depends on its initial mass. Thus we predict that

the most luminous helium stars will be the youngest.

58
as observed among the WN stars.

8. The most luminous helium stars will have

the largest radii (in accordance with our empirical results for WN stars) , but they do not appear

to have significantly higher effective temperatures.

This is fortunate since observationally we found

the excitation to be anticorrelated with the liaminosity

9. It is probable that the evolution is sen- sitive to the initial helium abundance. For example, the amount of matter left in the hydrogenrich envelope may be altered. Such a mechanism

may be responsible for the distribution differences

that I have interpreted in terms of initial chemical composition. Unfortunately, we cannot at this

time be more definite about this point.

10. It has often been suggested that the WR

components of binaries are filling their critical

Lagrangian lobes. If this is so, then stars at

point F appear the most likely candidates for WR

stars. Stars at point F have similar luminosities

but lower temperatures than stars at point E. Thus

the bolometric corrections would be less, the stars

would be brighter in the visual spectrum, and the

values of AM^ in Figure 6 would decrease in favor

of the helium stars. Point F also has the advantage of possible exotic atmospheric compositions

due to mixing induced by thermal pulses.

XI. DO THE WR STARS FORM A CLASS IN THE

SENSE DEFINED IN SECTION I?

The WR class was defined by common spectral

features and by the exclusion of planetary nuclei

and novae. The question now is: do we have any

evidence that the members of the class so defined

do not satisfy the criteria stated in Section I,

viz: (A) the stars are at closely related stages

of evolution, and (B) the principle mechanisms

responsible for the defining spectral characteristics are the same for all stars in the class.

First, consider criterion A. It seems a rea- sonable working hypothesis that the WR stars we know to be members of binaries have suffered mass exchange and have histories very like that sketched

in the preceding section. If this is so, are we

forced to assume that all WR stars are binaries?

Is mass exchange in close binaries the only mechanism that will produce a WR star? Kippenhahn (private communication) suggests that the answer is

59
"not necessarily". The basic property that makes

the helium star into a WR star must be its instability. If this instability arises from the temperature-sensitivity of helium or carbon burning,

then the specification for a WR star is that it has

a helium or carbon core with only a thin overlying

layer of hydrogen-rich material. A close binary is

a very convenient way of stripping off the outer

hydrogen layers to create such a configuration, but

there may well be other ways. However, no other

mode of envelope stripping is likely to be as fast, or begin as early in the evolution of the star, as that which occurs in a close binary. Hence the observation that binaries have a greater tendency to

concentrate to young associations than do single

stars, is a crucial one. As I have emphasized, the

observation is tentative, because the distances of

binaries are so uncertain. We badly need some way

to check this possibility.

In this connection, the nuclei of planetary

nebulae come to mind. Some of these have spectra

that are nearly identical with those of Population I WR stars. Planetary nebulae are believed to represent a post-red-giant stage of evolution. Are we,

here, looking at stars that have found a slow way

of reaching the critical configuration?

Of the WR subclasses, I think that the WN7

stars are the most likely candidates for non-membership in our hypothetical class of WR objects. The

spectra of these WN7 stars are very like those of

the Of stars. There is no reason to suppose that

Of stars are close binaries. Thus, if we decide

that all WR stars are close binaries, the WN7 stars

may well be excluded from the class.

We have no direct evidence on criterion (B)

However, spectroscopically similar objects that satisfy criterion (A) will most likely also satisfy

criterion (B) . The main use of the latter criterion is to clarify the relationship between the many

objects that are spectroscopically similar; that is,

a similar excitation mechanism may produce similar

spectra in objects that are unrelated in an evolutionary sense. Thus it is my personal opinion that

the planetary nuclei and single and binary WR stars

of Population I all satisfy criterion (B) , but must

be regarded as distinct classes if they result from

different evolutions and therefore do not satisfy

criterion (A) . It seems possible that the spectra

of Of stars and of novae result from different

sources of excitation and are not related to the

first mentioned varieties.

60
XII. SUMMARY

The correlations between the observed and deduced properties of WR stars and their spectral subclasses are shown schematically in Table 8. The arrows indicate the direction of increase of the

property within each sequence. A dashed arrow indicates uncertainty. Note that I have not arranged

the WN subclasses in numerical order. Inversion of

the order of the WN7 and WN8 stars makes the luminosity and age sequences more nearly monotonic.

TABLE 8 SUMMARY: OVERALL PROPERTIES OF WR STARS

WN3,4,5,6,8,7 Property WC5,6,7,8,9

.N lines

—> <c

^ 10 M©

WN6 > rest

Immediately postmain sequence

WN5 ,6,8

WN5,6,8

Excitation

Line /strength

|width

Luminosity

Size

Temperature

Mass

Age

Initial abundance of heavy

elements

Evolutionary

stage

Ring nebulae

Strong violet

absorption

edges

lines

<^

>

Cooler?

> 5 M^

WC5 < rest^

None

WC9

Note: Direction of arrow indicates increase.

61
(The exception is subclass WN3 in the liiminosity

sequence.) Considering only the nitrogen ions,

the excitation sequence is still monotonic after

this inversion of WN7 and WN8 , because the difference between these two classes depends on the

strengths of the He I lines.

To some extent the deduced properties depend

upon the assumption that we have a uniform class in

the sense discussed. If an entire subclass is re- moved from the class as a whole, the conclusions

drawn in this paper and given in Table 8 will not

be greatly affected. However, should we have a complete admixture of single and binary stars with

very different evolutions, ages, and possibly luminosities, then we should have to reconsider very

carefully some of our conclusions.

I have advocated two major interpretations of

these observations;

1. The classification system may be interpreted as a two-parameter sequence, initial mass

and initial chemical composition. The subclasses

of the WC sequence are determined primarily by

initial abundance differences among the heavy elements. The subclasses of the WN sequence are con- trolled primarily by the age (initial mass) of the

star. I realize this is an oversimplification.

2. A WR star is a helium- or carbon-burning

star, with only a very thin hydrogen-rich envelope.

These stars are pulsationally unstable, and this

instability generates shock waves that produce an extended atmosphere and an emission-line spectrum.

Mass exchange between the components of a close binary can lead to such a star; WR stars that

are members of binary systems have probably gone

through such a process. This mode of evolution can explain a major fraction of the observed correlations

between luminosity, age, etc.

Whether a single star can also attain a WR

configuration by some other process remains undetermined. Another such process would probably be

slower than mass exchange, and single WR stars are therefore expected to be older than binary WR stars.

Nuclei of planetary nebulae that display WR characteristics may represent the products of slow singlestar evolution to a WR phase.

We do not yet know the relationship between the

violence of the pulsational instability and other

properties of the star. This is probably the vital

link between the structural and evolutionary properties of the star and its atmospheric and spectral

characteristics

62
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Shajn, G. A. and Hase, W. F. 1953, Pub. Crimean Ap

Obs. 10, 152.

Shapley, J., and Nail, V. McK. 1948, Proc. American

Phil. Soc, 92., 310.

Smith, L. F. 1967a, P.A.S.P.^ 79., 347.

Smith, L. F. 1968a, M.N.R.A.S, , 138., 109.

Smith, L. F. 1968b, M.N.R.A.S.^ 140 . Smith, L. F. 1968c, M.N.R.A.S.^ in press.

Underhill, A. B. 1966 , Observatory ^ 86 , 25.

Underhill, A. B. 1968, Ann. Rev. Astr. and Ap . ^ 6., 39.

Van der Borght, R. , and Meggitt, S. 1963, Australian

J. Phys.^ 16., 68.

Van der Borght, R. 1968 (private communication)

Wallerstein, G. 1968 (private communication)

Westerlund, B. E. 1966 , Ap. J., 145., 724.

Westerlund, B. E. and Smith, L. F. 1964, M.N.R.A.S.y

128, 311.

Williams, J. A. 1966, A. J., 71, 615.

64
DISCUSSION

Chairman: Cecilia Payne -Gap osohkin

Payne-Gaposohkin : We have a number of copies

of the Hiltner and Schild paper on classification;

as these contain reproductions of spectra, I suggest they be distributed for your reference. [See

Part B, Figures 24-28, pp. 175 - 179.]

Now before launching into general arguments

on the material summarized, I suggest we call for

any additional facts not hitherto presented.

Steoher: The large majority of facts available

to us on Wolf-Rayet stars have come from observations in the visible spectral region. Observations

in the rocket ultraviolet can now provide us with

additional information both on the continuum and

on the line spectrum.

Our observations were made from an Aerobee

rocket carrying a thirteen inch f • 10 telescope with

an attached photoelectric spectral scanner with

three exit slits. The resolution was 10 A, and the

wavelength interval scanned was from 1150 to 4000 A.

The results were telemetered to the ground, recorded

on digital tape, and processed in a computer. The

whole rocket, with payload attached, was pointed to

the programmed stars with an accuracy of ten to

twenty seconds of arc, depending on the magnitude

of the star. This was accomplished with a startracker and a gyroscopic inertial platform. The

payload was parachuted to the ground and recovered

in good condition. Two scans of 20 seconds dura-

•tion were made on each programmed star. On this

flight we observed a Canis Majoris, e Canis Majoris,

C Puppis and y Velorum. The first two spectra appear to be interpretable in terms of standard atmospheric theory and will not be discussed here. C Puppis is an 05f star, and one component of y Velorum, Y2 Velorum, is a binary, WC8 + O.

Figure 7 is the spectrum of C Puppis in the

1800 to 3100^A range. The flux is in units of 10^

ergs cm"^s"^A~^at the top of the Earth's atmosphere.

This wavelength region is free of strong lines, so that by comparing the energy distribution with that

of a model atmosphere, we can obtain an effective

temperature. Two questions arise: (1) Is the model

sufficiently representative of the star to permit

us to integrate over the energy distribution in

65
ZETA PUP

1 —»

>•

——X

1 r— \

2000 * 2500 3000 WAVELENGTH X Figure 7, The spectrum of C Puppis in the 1800

to_3100 A range. The flux is in units of 10^ ergs

cm ^sec~^ A'^^at the top of the Earth's atmosphere.

unobserved spectral regions, especially below the

Lyman limit; i.e., is the physics sufficiently com- plete? (2) What is the correction for interstellar

extinction? Using the continuous models of Mihalas

(1965, i4p. J. Suppl.^ 9, 321.) with no correction

for interstellar extinction, we obtained an effective temperature of 2. 8x10 "^^K, which is of course ^ incorrect. The correction for extinction at 2000 A

is about 5 times the B-V color excess (Stecher,

1965, Ap. J,^ 142., 1683.). With the most probable

color excess, the effective temperature comes out

to be 4x10** °K. At 2000 k the difference between a 4x10 ''^K and a 5x10 '*°K model is only one tenth of a magnitude, when the models are normalized to the

visual flux. This is because at these temperatures,

2000 A is still way out at the red end of the Planck

curve. Thus an error of +0.02 mag in the color excess could increase the temperature from 4xlo'*°K to

5x10** °K. An error of this size is a distinct possibility. The color excess and temperature of C Puppis are important to our discussion, since y Velorum appears to have the same energy distribution in the 2000 A region. If we assume, as has

been suggested by several investigators, that C 66
Puppis and y Velorum are physically associated in

space, then the correction for interstellar extinction will be the same, and the effective temperatures will both be about 4x10** °K.

Figure 8 shows the short wavelength region of

C Puppis. The remarkable thing to note is the P Cygni type profile in the resonance lines of C IV,

Si IV, N V, C III, etc. The excitation of this

star is sufficiently high that the N IV line at

A1718 is also in emission with a Doppler-shifted

absorption edge. The level is 16 volts up and

should be useful in determining the structure of

the envelope. The wavelength difference between

the emission peak and J.he blue-shifted absorption

minimum in C IV is 18 A, which corresponds to a ve- locity of 3400 km/sec. While this is an upper

limit, the real velocity, which depends on the

model, will certainly be large. The time constant

on the amplifier was set to rise to 1/e of the true

ZETA PUP

25

20

15

10

cm

cm

cnr

Si 17 Hen

; s

H I

\

• Hen

Si IV

1 ... __i - / 1 1000 1500 2000 WAVELENGTH X

2500 Figure 8. The spectrum of C Puppis in the 1200

to the 2400 A range. The flux is in units of 10^

ergs cm" sec sphere

at the top of the Earth's atmo67
signal level over the scanning bandpass of 10 &. When corrected for instriimental profile, the center of the absorption line should be quite black,

indicating a large steady-state mass loss.

The absorption at X1216 is due to interstellar

hydrogen. The line is formed by radiation damping

and can therefore be used to obtain the amount of

neutral hydrogen between the Earth and C Puppis.

The small equivalent width, 4 A, indicates a low

column density for neutral hydrogen, consistent with

the small reddening correction in the visual region.

Finally, one should note that in many respects the

ultraviolet spectrum of this star is similar to that

of a Wolf-Rayet star. In fact, if we were to ignore

the wavelength region, this spectrum would meet most

of the criteria so far discussed for inclusion among

the Wolf-Rayet phenomena (and wavelength region was not one of the criteria)

Having presented the "control star", we proceed to Figure 9 which is the 1800 to 3100 A scan

GAMMA VEL

!

i

:•

1 [cm]

1

i

\ !

c in;

1

!

1

i

1

' —

:

1

... —

i

i

1 i i;

25

20

15

10

2000 2500 WAVELENGTH X

3000 Figure 9,^ The spectrxim of y Velorum in the

1800 to 3100 A r^nqe. The flux is in units of 10^

ergs cm"^ sec"^ A" at the top of the Earth's atmosphere.

68
of Y Velorum. The continuum is less obvious, but

a reasonable estimate can still be made. In addition to the Wolf-Rayet continuiam there is the 0-

star and a small contribution from yi Velorum. The

lines of the 0-star are not noticeable, so it is

reasonable to assume that most of the continuum is

due to the Wolf-Rayet component itself. With these

assumptions , and the same correction for extinction

as we applied to C Puppis , the effective temperature is 4xio'*°K. This is of course assuming that a hydrogen-helium atmosphere can be applied to this

star. If a he liiim- carbon atmosphere should apply,

changes in the far-ultraviolet could alter the temperature. A check against Hanbury Brown's angular

diameter, the observable flux, and the flux of the

model at X190 0 showed good agreement.

The strong line at A2296 is the strong permitted to transition in the C III singlets.

The even more intense line at X1909 is the intercombination line of C III. This line has an os- cillator strength that is smaller by a factor of a million than that of the permitted line and thus

poses the interesting question of the processes

of formation. West and I have been looking at this

problem, and he will discuss our calculations later.

It appears that stimulated emission is necessary

to produce the C III line. The required overpopulation of the triplets with respect to the singlets

is most likely due to the dilute radiation field in

the envelope (Struve and Wurm, 1938, Ay. J. ^ 88 , 84.) and/or to mechanisms as yet unspecified. The

result is laser action.

The only other line clearly identified in the

spectrum is due to C IV at X2524. The atomic data

for the high excitation present in these stars are not sufficiently complete to allow satisfactory

identification of most of these lines. It should

be remembered that the lines in the ultraviolet con- tain a considerable amount of energy; the continuum

is many times more intense in the UV, and if a line

is to show up, it must be intense.

Figure 10 shows the far-ultraviolet spectral

scan of y Velorum. It overlaps the previous spectrum and again shows the X2296 and X1909 lines of

C III. The Ha line of He II appears to be quite

broad. The broad general absorption between A 1700

and A1300 is due to molecular oxygen in the Earth's

atmosphere. There is sufficient information to

correct for it, and this is now being done. The

C IV line at X1550 appears both in emission and in

absorption, again indicating an expanding envelope.

69
GAMMA VEL

25

20

10

5

0.

n TTT ' r» TTT I L HI L 111

j _ . ;. 1 . 4

[cn H n '

i 0 i

i _

1

1 c n

'

. : He n

V; O TTT L 111 - cn

i

' _ _

i

i

•

': : y

• '

t _

Hetn Si

1

!

]

i

• »v'

\

- ^A

ciy

A— ^

i

' Si m

\ —.—

1 / 1 'lOOO 1500 2000 2500 WAVELENGTH X Figure 10. The spectrum of y Velorum in the

1200 to 2400 A range. The flux is in units of 10^

ergs cm~^ sec""^ A-^at the top of the Earth's atmosphere . although the velocity of 1500 km/sec is less by

a factor of two than the velocity in C Puppis

This is also evident in the resonance lines of C II

,

Si IV, etc. Again a large number of lines belonging to the star are unidentified.

One interesting possibility for unidentified

lines is the question of the physical reality of

the quark. This mathematical particle which has

had such great success in SU3 theory may or may not

exist in nature. If quarks are physically real

they would presumably be produced through cosmic

ray reactions in the interstellar medium, would be

thermalized, and would be present through star formation in younger stars. At interior temperatures,

the quark of charge - 1/3 would bind itself to C,

N, and O. The Wolf-Rayet stars offer the most

likely possibility of detecting them by observing

the spectral lines of carbon of charge 1+2/3,

2+2/3 and of nitrogen of charge 1+2/3, 2+2/3, etc.

The long path length through the envelope of the

star, coupled with the probable overabundance of

70
these elements, can greatly increase the current

upper limit on these particles. Accurate theoretical wavelengths are needed for the quarked atoms

,

along with more complete multiplet tables.

Nitrogen V appears in absorption in y Velorum,

and N IV is probably present in emission. The

lines are weaker than in C Puppis but are sufficiently clear to establish the presence of nitrogen

in a carbon Wolf-Rayet star. A number of interstellar lines appear in the spectra of both stars:

C II, Si II, Al II, O I, Fe II, Ce II. There is also

the possibility of autoionization of Al I and Ca I. The equivalent width of hydrogen Lyman-a in y Velorum is similar to that in C Puppis, indicating

a low column density of neutral hydrogen and implying the correctness of the interstellar-extinction correction.

Undevhill : I have quite a few comments, but

I won't make them all at once. First I would like

to report on yz Velorum. Ganesh and Bappu (1967,

Kodaikanal Bull.^ 16, No. 183.) have made a radialvelocity study of this binary and find a period of

78.5 days. They determined velocity curves for

the Wolf-Rayet component from three different lines,

and they present arguments to show that the mass

ratio, M^r/Mq = 13/46, derived from the C III com- plex at 4652 A is probably the best value. They

also conclude that sin i is about one, or at least

that it is greater than 0.8.

Now the spectral class of y2 Velorum is, ac- cording to Lindsey Smith, WC8 + 07, and her estimates of the absolute magnitudes of the subclasses

WC8 and WC9 are based entirely on this system.

Using Graham's distance for yi Velorum, (derived

from H3 and UBV photometry) , she came up with an absolute magnitude of -6.6 for the system and -6.2

for the WC8 star. Now Ganesh and Bappu note that

yi Velorum (HD68243) and ya Velorum (HD68273) are

an optical double, that yi Velorum is of MK type

B2IV, and that it is 2.4 mag fainter than Velorum. If you assume that yi Velorum has the standard absolute magnitude of -3.3 for a B2IV star,

and that the absorption and reddening correction

are the same for yi and y2 Velorum, then the absolute magnitude of the system, y 2 Velorum, is -5.7.

If you assume further that the two components of

y2 Velorum are equally bright, then the WC8 star

has an absolute magnitude of -5, which differs by

1.2 mag from Lindsey 's value.

Sohild: How did they get the classification

B2IV for yi Velorum?

Underhill : They say that an absolute magni71
tude of -3.3 is ascribed on the basis of the MK

classification type assigned. They did not give

a source for the classification; I suspect they

took it out of the Bright Star Catalogue.

Payne-Gaposchkin : Does the Bright Star Catalogue give a reliable luminosity class?

Sehild: Satisfactory MK classifications are available for some of the brightest B-stars, but

not for all of the southern OB-stars. This is a crucial point, because you are basing an absolute

magnitude estimate on a single spectral and luminosity classification.

Underhitl : Well, I just wanted to present

this difference in results. Personally I prefer

the fainter magnitude because, with the exception

of WN7 and WN8 stars which have quite different

spectra, most WR stars have absolute magnitudes of

the order of -5. The adopted absolute magnitude

will certainly make quite a difference to the distribution of WC8 and WC9 stars.

In this connection, I would like to refer to

Campbell's hydrogen envelope star, BD+30°36 39 or HD184738. It was the original WC8 star, but over

the course of years it has been placed with the

planetary nebulae. Its apparent visual magnitude

is 10.0, and if you accept Seaton's distance, its

absolute visual magnitude is -4.9. Admittedly,

Seaton warns that his method for determining distances is very uncertain, particularly when the

nebula is as dense as this one; but if you want the

star to be as bright as -6, you will have to put it

a long way off, and then I ask how you are going to see such a small nebula.

Smith: On the subject of the absolute magnitude of Y2 Velorum: I personally would regard a luminosity for yi Velorum determined from the H3

photometry of Graham as more reliable than a luminosity derived from a spectral classification of

unknown source. Regarding the absolute magnitude of the WC9

stars: If you drop the absolute magnitude much

fainter than -6.2, then you must find another ex- planation for their peculiar angular distribution.

Confinement within 45° of = 0° implies confinement within 7 kpc of the galactic center (see Figure 1). If the luminosity I gave is correct, then

nearly all WC9 stars are sufficiently far away from

the Sun to be within 7 kpc of the galactic center.

If, on the other hand, the luminosity is fainter

than -6.2, then the WC9 stars are closer to the Sun

and further than 7 kpc from the galactic center.

72
You must then explain why they are not found in

directions greater than 45° from i^-^ '= 0°.

Underhill : Suppose we are venturesome and

agree that WC9 stars are rather older than some of

the others. Am I correct in saying that on the

whole the stars toward the center of the Galaxy

are a bit older than those in the outermost arms such as the arms in the anti-center region of Orion?

We are certainly not talking about stars that are very old. Well, the fact that the 6 or 10 WC9 stars

that we do observe just happen to be in that quadrant doesn't bother me much: toward the center of

the Galaxy you have more arms in which to see them,

and after all, the Poisson distribution allows you

an uncertainty of the square root of 10 which is 3.

So the observed distribution could be just chance.

Finally, I think it is essential to draw in the

dark, obscuring lanes before concluding anything

about the distribution of objects at distances

greater than 3 kpc from the Sun.

Smith: But I do see WC7 and WN6 stars more

than 45° from = 0° and at distances comparable

to those of WC9 stars.

Consider Anne Underhill 's second point, the

similarity between Campbell's hydrogen envelope

star and the WC9 stars. In the southern hemisphere

there is a second, similar object, Henize 99, with

a WC9 spectrum plus nebular lines. A spectrogram

has been obtained by Louise Webster. Whether these

objects are genuine planetary nebulae remains to be

determined. However, I have three reasons for as- serting that they are different from the stars that

we call classical WC9 stars. (1) Stellar emission

lines in the spectra of Campbell's star and of H99

are consistently narrower by a factor of about 1.4

than the emission lines in the spectra of classical

WC9 stars, i.e., they are about 6 A as against

about 8 A in WC9 stars. (2) We do not see nebulae

around any classical WC9 stars, although three of

them have apparent magnitudes brighter than either

Campbell's star or H99; nor do we see nebular lines

in the spectra of any classical WC9 stars. Thus I assert that the classical WC9 stars do not have

nebulae. (We do see nebulae around WR stars, but

as I mentioned earlier, these are restricted to

three subclasses within the nitrogen sequence.) (3)

Whereas all classical WC9 stars are concentrated

within 45° of the galactic center, the two stars

with nebulae lie in directions outside this range.

Furthermore, I think that the distance of Campbell's

star is extremely uncertain and that an absolute

73
magnitude based on it is no better than an estimate

based on Y2 Velorum.

Payne-Gaposohkin : You said there was a suggestion of spiral structure in the distribution of

galactic WR stars. I wonder if you could show us how it runs. Smith: There are three nearly circular arms located 10, 8 and 6 kpc from the galactic center.

The best fit is obtained with arms inclined about

5° to the tangent direction, in the sense that the arms are trailing (see Smith 1968c)

Payne-Gaposohkin : But these must be "classical" spiral arms, not arms defined by the Wolf-Rayet

stars

Thomas: I am concerned with the distinction

between a WR object and the possibly broader range

of phenomena which exhibit WR spectra. Lindsey has

concentrated on WR objects; she asserts that there

are indeed such objects. Let me summarize my own understanding of her picture and ask her to correct

and elaborate it. A WR object starts out with a definite chemical composition and mass, and possibly with a definite location relative to other

objects. During the course of its evolution, it

reaches the WR stage, i.e., its spectrum shows the

characteristics which define the WR class. Lindsey

asserts that at this stage the star consists of a helium core surrounded by a thin hydrogen shell and

that it is overluminous for its mass. The WR spectrum results from the instability of this physical

configuration: radial pulsations produce shock waves

that mechanically heat the atmosphere.

So there are two questions: (1) How do you

produce this physical configuration from the initial

conditions in the stars? Is there a certain range

in initial mass and chemical composition for which

a star is guaranteed to pass through the WR stage?

(2) How does this physical configuration act to produce the WR spectrum? What is the special relation

between mechanical energy, momentum, and chemical

composition that produces the characteristic spectrum? Now you permit variations in spectra within

the WR class, namely those variations that corre- spond to the different subclasses. So you would permit variations in the initial configuration of the

star and in the mechanism producing the spectrum.

You began by suggesting that the subclasses corre- spond either to different positions along a single

evolutionary track, i.e., to different stages in the

evolution of stars of the same initial mass and

chemical composition, or to similar positions on dif74
ferent evolutionary tracks, corresponding to differences in initial mass or chemical composition. It

is my understanding that you conclude by adopting

the second alternative, since you suggest that the

WC subclasses correspond to differences in initial

chemical composition, and WN subclasses correspond

to differences in initial mass. Now in answer to

(1) , you state that a close binary system will cer- tainly produce the necessary configuration, and

without being specific you suggest that there may

be other alternatives for a single star. But you

also state that the helium-core hydrogen-shell con- figuration may not be the only way of mechanically

heating the atmosphere. So it is not obvious to me that you are claiming . that the WR objects are the

only class of objects that uniquely produce the WR

spectrum.

Smith: I made the hypothesis of a uniform

class and showed the arguments for and against it;

I do not necessarily believe the hypothesis is

correct. I think that to produce a WR spectrum,

you need the mechanism which I discussed and that

it does not matter whether the star is single, binary, novae or whatever. However to obtain a uniform class, we require that all the stars be at

the same evolutionary stage. It is at this point

that we have to ask very carefully whether they

are binary, single, novae, or nuclei of planetary

nebulae. If a star reaches a certain configuration of mass, core size and composition, then it

will be, say, a WN7 regardless of how it got there.

But if you have different modes of evolution, you

may have differences in distribution. For example,

single stars are likely to be older, and this will

be reflected in their distribution within the Galaxy. Again, you may have different distributions

cimong the subclasses. For example, the central

stars of planetary nebulae may represent single

stars which have achieved the WR configuration.

Since most planetary nuclei are WC stars, it may

somehow be easier for a single star to arrive at

the WC configuration.

Thomas: In your answers to (1) and (2), you

have in no way differentiated between WC and WN

stars, either in terms of the model or of the ex- citation mechanism. You have only specified a helium core surrounded by a hydrogen shell. For all

you have said, the difference between the WC and

the WN sequences could be just a difference in ex- citation.

Smith: Yes, that is possible. However, it

75
does appear from the distributions that the WC

stars do have higher initial heavy-element abundances (including helium as a heavy element) . The

most likely suggestion is that the initial heavyelement abundance affects the evolution. In particular it may affect the amount of hydrogen-rich

material which remains around the helium core after mass exchange. This, in turn, will almost cer- tainly affect the instability properties and hence

will affect the amount of mechanical energy fed

into the outer atmosphere. It may also affect the

surface composition: once you pass the helium burning stage, you have two thin shell sources, one burning helium, the other hydrogen. These are likely

to produce thermal pulses and mixing in the outer

layers. The energy content of the thermal pulses

and the degree of mixing to the surface will depend

on the thickness of the hydrogen-rich envelope,

hence on the initial heavy element abundance.

Thomas: You are bringing more carbon up from

the interior by differential mixing. Is this the

kind of thing Paczynski suggested?

Smith: Yes, it's the kind of thing suggested

by Paczynski, although the models that go that far

and the detailed suggestions are due to Kippenhahn.

There may also be other mechanisms.

Underhill : 1 have two comments: one is on the

assignment of objects to the WR class; the other is on abundance differences. On Beals' first figure,

he pointed out one little spot we could not see,

under which was written "Oa" . He explained that in

the early days of spectral classification, these

objects were found on objective prism plates, that

they were very peculiar, and that there were only a few of them. Now they are called WC and WN, but

the point is that we identify each of these stars

by its spectrum in a very limited wavelength range,

between 4000 and 5000 A. Indeed if Stecher had

shown you the spectrum of C Puppis only in the

range X < 2000 A, you would immediately have said

it was a WR star, whereas we all know that in our standard classification system, it's an 05f. So

there is no guarantee that we haven't got a mixture

of widely differing objects in our WR class. Investigations such as Lindsey's are aimed at trying

to separate these objects into smaller groups. But

let's not insist on too tight a spectral identification, at least until later in the symposium when we hope to be able to answer a few more questions.

Now regarding abundance differences: Those

stars which Lindsey suggests are formed toward the

76
center of the Galaxy are supposed to have more heavy

elements than the so-called "normal" population in

the solar neighborhood. They are supposed to be

old stars. Now to support her arguments relating

the abundance of heavy elements to age and galactic

distribution, Lindsey has quoted work on the cepheids.

There are supposed to be both strong- and weak-line

cepheids, and it is often suggested that strong

lines imply a greater abundance of heavy elements.

And so they do, on the normal LTE approach. Now

I've been trying to interpret stellar spectra by

standard textbook theories for 20 years, not because I believed the textbooks, but because there

wasn't anything else. In 20 years I've computed a lot of models and a lot of lines. Superficially

the models look like stars, but when you come to

compare them in detail, they aren't stars. And

when I look into the physics of what I am doing,

and when I talk to the plasma physicists, I find

I've been using the wrong theory and getting the

wrong answers . For example , the LTE approach leads me to conclude that many more atoms contribute to

the formation of a strong line than would be re- quired if the correct theory were used to interpret

the spectrum. There is also strong evidence that

any star with a slightly extended atmosphere will

have broad lines due to motions. So if you say

strong lines indicate greater abundances, you are wrong. So I think you are on weak ground in saying there is a definite anomaly in the abundances,

either of heavy elements or of helium, in the WR

stars

Thomas: Anne, you know I'm the last person in

the world to stifle criticism of the conventional

methods of atmospheric analysis, but concerning

abundance differences, I think it is essential to

differentiate between arguments based on atmospheric

analysis and those based on interior models. It seems to me that all Lindsey 's arguments are based

on the latter.

Payne-Gaposohkin : Some of the arguments about

cepheids suggest that their evolutionary paths differ and that this indicates differences in composition.

Smith: Basically the argument is this: we have gross variations of population properties in

different galaxies and in different parts of the

same galaxy; nothing else will account for these

variations as well as the suggestion that they re- sult from differences in initial chemical composition. The fact that the atmospheres also appear to

77
vary in composition provides additional verification of this hypothesis.

Thomas : I thought your argument was based on computations by the interior people , Kippenhahn and

Paczynski; if you are now turning to atmospheric

analyses, I'll back Anne's objections.

Smith: Mine is an interiors argument. Properties of the cepheid models depend on their initial

chemical composition; see, for example, Hoffmeister '

models in Zs. f. Ap . , 65^. She has calculated two

sets of evolutionary tracks with different Y's and

Z's; it is impossible to tell whether Y or Z is the

sensitive parameter. I'm suggesting that the evo- lution of a WR star is affected in a similar way by

its initial chemical composition. Mine is therefore an interiors argument.

Payne-Gaposohkin : Hof fmeister has the extreme

and the ordinary Population I in her models , but

they were very similar; while the one with low Y was very different. Well, we'll argue about it

another time.

Underhill : I am still worried about the helium

abundance. Observations by Bappu indicate that there

may be H mixed with He II in the Balmer series of

some of the well-known Cygnus stars. The fact that

you cannot easily observe H is obvious from estimates

of electron temperature in these atmospheres. Your

gas won't be cool enough to produce strong Balmer

lines until you're so far out in the atmosphere you

won't see it against the star anyway. So the interior question is wide open; you have nothing

against which to check your model. If you're going

to start off with a large helium and low hydrogen

abundance, the atmosphere has to reflect this. I just don't think there is enough observational evidence to support that position. It may be perfectly

true that by starting with a close binary, you can get mass exchange, but how do you know that you don't

end up with two 0-stars? There is not one bit of

evidence that either of the stars will have a WolfRayet spectrum. So the first thing we have to do

tomorrow and the next day is to answer the question:

How do we produce a WR spectrum?

S. Gaposchkin : I would like to emphasize that

V444 Cygni is an almost perfectly determined binary

system, both spectroscopically and photometrically.

One must therefore accept the values, mass = 9 . 5 M®

and luminosity = -2.8, as well established. So

Miss Smith has taken the WR star too bright for this

particular subclass.

Smith: The absolute magnitude for the WR com78
ponent depends on an assumed luminosity ratio.

5. Gaposahkin : False! It is an observed ratio!

Payne-Gaposohkin : We will defer further discussion until after Kuhi has discussed this star

tomorrow.

Banbury Brown: I am going to show you some measurements made on Y2 Velorum with the stellar

interferometer at Narrabri observatory in Australia

(see Figure 11) . This instrument is an intensity

interferometer, and it has the property that the

signal to noise ratio is independent of the optical

bandwidth, provided only that the optical bandwidth

is large compared with the electrical bandwidth.

We have exploited this peculiar property to make

measurements of the angular diameter of the star in

the continuum and in the light of an emission line.

The interferometer measures correlation as a function of separation between the two mirrors, or baseline. It can be shown that this correlation is

proportional to the square of the fringe visibiliFigure 11. The interferometer at Narrabri

Observatory, Australia. Each reflector, 6.5 meters

in diameter, is formed by a mosaic of 252 hexagonal

glass mirrors approximately 38 cm across and 2 cm

thick. The mirrors are not figured to a high precision because it is not necessary to form a con- ventional image.

79
ty observed with a classical Michelson interferometer with the same baseline. It follows that

from a curve of correlation versus baseline one can readily find the angular diameter of a star,

W§ made the measurements in the continuum at

4430 A with a filter bandwidth of ±50 A, and also

in the C III/IV emission line with a filter cen- tered on 4656 A with a bandwidth of ±12.5 K. The

center wavelength of the latter filter is displaced

from the center of the line to allow for the fact

that due to aberrations in the optical system, much

of the light does not transverse the filter normally. Figure 12 shows the measured values of

correlation as a function of the baseline in meters.

The error bars on each point represent the uncertainty due to statistical fluctuation of the correlator output. The interpretation of the continuum

results (upper part of figure) is complicated by

the fact that we are observing a binary star. The

correlation from a binary is a function of the angular separation, position angle and relative brightness of the two components, and unless there is some auxiliary optical data, especially about the

period, it is very difficult to interpret the re0.5

MARCH -APRIL 1958

72 VEL 4430 & ±5oX

20

0.5

13

: {

60 100 METERS200

^2 VEL 4656 X±I2.5&

J Li . L 20 60 100 200

METERS

Figure 12. Correlation observed from yz Ve- lorum in the continuum (upper figure) and in an emission line (lower figure)

80
suits. When the baseline of the interferometer is

very short and the angular separation of the binaryis not resolved, the measured correlation is proportional to (Li + Lz)^, where Li and L2 are the

luminosities of the two components. On the other

hand, when the baseline is long and the angular

separation of the binary is resolved, the correlation is proportional to (Li^+ Lz^) and may therefore be considerably less. In the transitional re- gion between these two conditions, we get all sorts

of complicated results which depend in detail on the

spacing and position angle of the binary and also

on the baseline.

Looking at the results in the continuum, we see

at once that we need to reduce the errors by making

longer observations; nevertheless, there are some things which can be deduced from them. First we

can derive the angular diameter of the star by fitting a theoretical curve (for a uniform disk) to

the three points at the longest baselines; we can then repeat this, including the four and five longest baselines, and we can compare the results. We

can be reasonably sure that the three, and perhaps

four, longest baselines are so long that the separation of the binary will be resolved. At the two

shorter baselines we cannot be sure that the binary

is resolved, and you can see from the figure that

the correlation at the shortest baseline is significantly above a smooth curve through the other

points; this suggests that, at the shortest baseline, the binary was not completely resolved, and

I shall assume that it was not resolved until the

baseline exceeded 50 m. The angular diameter of

the brighter component, deduced from the four

longest baselines is:

0 (continuum) = 0.44 ± 0.07x10 seconds of arc.

I hope you will appreciate that this is a preliminary result, as the observational data were taken

only a few weeks ago and have not yet been completely

analyzed. From the intercept of the fitted curve with the zero-baseline ordinate, we can find the

ratio of the luminosities of the two components of

the binary. Our preliminary result is:

Am = 1.3 ± 0.6 mag

We cannot, of course, tell from the measurements

which of the two stars is the brighter; but roughly

speaking, if they differ significantly in bright81
ness, as they appear to do, then the interferometer yields the angular diameter of the brighter

star. Thus if the WR component is brighter than

the 0-star, its angular diameter is given by the

value above.

The lower part of Figure 12 shows the corre- lation observed in the emission line. I should

have mentioned before that the correlation is

normalized to the value expected from a point

source giving the same light flux. Thus the mea- surements in the emission line point to a single

source with an angular diameter of

0 (emission line) = 2.1 ± 0.3x10"^ seconds of arc. Again I must emphasize that the result is preliminary. For example, one must subtract from the

"emission line" results the contribution due to

the continuum, and this has not been done properly

yet. I have not yet had time to do it. Nevertheless, I do not think that the two values of angular diameter which I have quoted will be changed

very much by a more complete analysis.

From the results one can see that, assuming

that the WR star is the brighter of the two, its

apparent diameter in the emission line is about

4.8 times its diameter in the continuum. Also, if we take the value of its parallax given by Allen,

160 pc, then the radius of the WR star is roughly

7.5 Rq, and that of the emission region is roughly

36 Rq. Combining the measurements of angular size

with measured values of monochromatic flux (fx)

,

we find, again very roughly, that the brightness

temperature of the WR star at 4430 A is about

3. 1x10 '^^K, and the brightness temperature of the

emission region is about 1.2xlO'*°K. A more detailed analysis will probably alter these values

a little, but I don't think very much.

S. Gaposchkin : Do you have a parallax?

Hanbury Brown: The value of parallax, 160 pc,

which I used was taken from Allen's "Astrophysical

Quantities", a standard work. I regret that I do

not know the original source. I should like to

add that one is seriously handicapped in interpreting the continuum observations on this star by

lack of spectroscopic data. We badly need to know

its period. Miss Underhill has kindly drawn my

attention to the paper by Ganesh and Bappu which

gives the period as 7 8.5 days. This measurement

could be a great help to us in interpreting our

data.

82
Smith: Graham's distance of 460 pc, determined

by H3 and UBV photometry, is rather larger than

Allen's figure of 160 pc , so I would multiply all

Hanbury Brown's dimensions by a factor of about

2.5. The luminosity of the system is then -6.6;

for the 07 companion, I took Mv = -5.2 (SchmidtKaler 1965) which leaves -6.2 for the WC8 star.

Thus the WR star appears to be the more luminous

by about 1 mag, which compares well with the fig- ure of 1.3 ± 0.6 from the interferometric observations . Payne-Gaposchkin : Can't you get some corroboration from the striking prominence of the bright

lines in the combined spectrum?

Smith: Yes, just because the emission spectrum

is so prominent, we would not expect the WR star to

be as much as 1.3 mag fainter than its companion.

But if the distance is indeed 160 pc, the luminosity

of the system is approximately -4.2, which is fainter

than an 07 star alone. I do not think this is likely.

Thomas: In your summary you suggested a value

of 5.0xl0'*°K for the photospheric temperature of WR

stars. Stecher quoted a similar figure for y Velorum as a whole. Hanbury Brown's figure is about half

of this. Comment?

Hanbury Brown: There is no reddening correction in our determination, and I don't know how

large it might be.

Underhill : If you assume that Y2 Velorum is at

Allen's distance of 160 pc - which is where Ganesh

and Bappu place it - and that the two components are equally bright, you get an absolute magnitude of

-5 for each of them. If you then use Hanbury Brown's

figure of 3x10** °K as an effective temperature, you

get a radius in agreement with his results , which

are consistent with the idea that a WR star is

rather like an 0-star.

Thomas: You mean in the continuum and in the

atmospheric regions where, the continuum is formed?

Underhill: Yes. As far as the continuum is

concerned, a WC star appears rather like an 0-star.

Of course this statement could be challenged by

Kuhi tomorrow.

Thomas: So the point at issue is whether the

continuum of a WR star behaves roughly like a blackbody at 5x10 ''^K, or like a blackbody at half that

temperature, or at a greater temperature, or not

like a blackbody at all.

Underhill: Certainly not much higher than

5xl0'*°K.

83
Payne-Gaposohkin : Anne, I would like to return

to a point that you brought up. You don't seem to

think the evidence from the cepheids concerning

differences in chemical composition is very con- vincing. How do you feel about the evidence presented by Preston on RR Lyrae stars?

Underhill : I haven't got the subject at my

fingertips, so I would rather not comment.

S child: Van der Bergh has been studying globular clusters both close to the galactic center and

further out. He finds a tendency toward metal en- richment in the stars of those clusters which are close to the center. I can't comment further, as

I only heard this reported at a Cal Tech colloquium.

Payne-Gaposchkin : W. W. Morgan's spectra are the best and most convincing evidence of this effect.

Underhill : The basis of all this work is nar- row-band photometry. Correlations are sought between color indices and observed spectra, some of

which show weak, some strong, lines. The color

indices are Stromgren colors, which give relative

brightness between say X4400 and X4200. Another

color is used to select objects that are similar

in the longer wavelength region and therefore have

the same effective temperature. The blue color index is then correlated with line strength. They

presume that because a line is stronger in one set

of stars, there is an abundance difference. This

is where I rise in wrath. True, in the standard

theory the abundance is your only free parameter;

I just don't believe the standard theory.

Sohild: I'm out of my field here, but I seem

to recall that Van der Bergh 's work consisted of a complete analysis of HR diagrams for the globular

clusters, and that correlations were observed between metallic abundances and the numbers of stars on two horizontal branches. Too bad Van der Bergh

is not here to defend himself.

Underhill : His work is almost exactly the same

as I described. He plots an HR diagram: b-v narrow- band-color index against apparent magnitude, which

is absolute magnitude ' because in a cluster all the

stars are at the same distance. He gets all the

stars lying in a certain sequence, and using model

atmospheres or the interpretation of the spectra of

a few bright stars, he says one turn-off point cor- responds to strong-line stars and another to weakline stars. Having established the relationship for one or two stars in nearby clusters, he then applies

it to fainter field stars. The technique is based

84
on an empirical correlation based on a rough interpretation of spectra using rudimentary theory.

Fayne-Gaposohkin : Preston's work is not based

on clusters, but on bright RR Lyrae stars. How do

you evaluate the work of Oke and his collaborators

on weak- and strong-line RR Lyrae stars and cepheids,

in which you get quite different relations between

Tgff and B-V for the different stars?

Underhill : Well, down to spectral type A5, we

have photometric spectral types and we have absorption-line spectral types. For the photometric types, we have a relation between the slope of the contin- uum and Teff established by model atmosphere calculations. If you go cooler than A5, log g becomes

the dominant factor in determining the slope of the

continuum. I think Preston is a good observer and

his interpretation gives me confidence. But although continuum spectral types do seem very con- sistent, there are still people like me who are old-fashioned enough to prefer absorption-line

spectral types.

Payne-Gaposchkin : And of course Preston uses

the difference between the two kinds of spectral

types as a measure of the metal abundance in RR

Lyrae stars. I must say that while I am not con- vinced about the cepheids, I do find Preston's work

pretty convincing.

Underhill: Basically the reason we are all

getting so excited about weak and strong heliumline B-stars is that the photometric spectral types

and the absorption-line spectral types do not cor- relate in a unique manner. I hesitate to interpret

this in terms of abundance differences.

Kuhi: Anne, would you believe abundance differences derived from photoelectric measurements,

independently of model atmosphere calculations?

These are based on conspicuous differences in the

strengths of lines in stars of a given temperature,

as defined by narrow-band indices free from absorption lines. I am referring to the work of

Spinrad.

Underhill : What I'm saying is that all deductions from photoelectric indices based on an empirical spectroscopic interpretation of the continuum

cannot be uniquely related to the absorption-line

spectra. I agree that large differences in the line

strengths are possible; the question is how to interpret them.

Thomas: Let me clarify this: First, Anne is

questioning not the observations but the physical

interpretation. It's not a matter of whether you

85
believe photoelectric versus spectrophotometric

results, but rather of whether you believe that a difference in line strength necessarily implies a difference in abundance. On the classical LTE

theory, it does. I agree with Anne; it is not obvious this theory is applicable, so it is not obvious the interpretation in terms of abundance

differences is valid. The second question is whether

stars with the same continuum temperature are nec- essarily similar in atmospheric regions above where

the continuum is formed. Your continuum temperature refers only to conditions at the place where

the continuum is formed. To assume that these con- ditions fix uniquely the conditions in the region

of line formation is a big assumption. If, for

example, the continuum is formed in a region where

conditions are controlled by the radiation temperature, and if the line spectrum is formed in a solartype chromosphere-corona, then the continuum and the

lines are formed under wholly different conditions

and wholly different mechanisms. There is no obvious reason why there should be any correlation

between them. Maybe there is, but this must be

proven, not assumed. I think this is the basis of

Anne's argument.

Underhill : Thanks Dick. I agree; you've put

it very clearly.

Thomas : But I think we have drifted from the

main point. Lindsey, you find that the different

subclasses of Wolf-Rayet stars have different distributions both within the Galaxy and among extragalactic objects. You then argue along the lines

used by Shapley many years ago when he found that

the mean periods of cepheids in the central regions

of galaxies differed from those in the outer re- gions. He said, "Maybe this is an indication of

gravitational potential or something". The current

approach is to interpret it in terms of abundance

differences. But this is only one of many approaches.

The important observational fact is not that the

abundances vary from one part of a galaxy to another,

but that some property varies and that this property

in some way affects the subclasses of WR stars, just

as it affects the mean periods of cepheids or of

subclasses of cepheids. Ever since Shapley, people

have been looking for this property.

You go on to note that theoretical calculations

show that differences in initial chemical composition could account for the differences in the periods

of cepheids. So you suggest that we might account

for the differences in the WR subclasses in the same

86
way (including, possibly, a difference in initial

mass). You stick to this suggestion, because it's

the only logical one you can see. Anne, you are not worried about Lindsey's arguments as such; you are questioning £he simple interpretation of differential line strengths in terms

of abundance differences , and you are worried be- cause Lindsey's arguments appear to support this

interpretation. I share your worries, but if we

accept that certain subclasses of WR stars and ce- pheids are concentrated toward the galactic center,

we have either to accept differential abundances or

to propose an alternative interpretation.

Smith: I agree with that summary. What has

happened over the years is that when we have looked

at stars in the halo and at stars in the disk, we

have found correlations between line strengths

,

dynamical orbits, and photometric criteria. We have

come to the conclusion that stars in the halo are intrinsically different from stars in the disk. By

far the most likely and logical interpretation is

that these differences reflect differences in initial chemical composition. It is my impression

that this interpretation is generally accepted,

i.e. , that it is believed that Population II (halo)

stars have lower metal abundances than Population

I (disk) stars. Now we have also come to think of

Population II as being old. In the center of the

Galaxy we find stars which we assign to Population

II because they are old; we have assumed that they

are also underabundant in heavy elements. But the

standard criteria by which we estimate metal abundances indicate that the metal abundance in Population II stars in the center of the Galaxy is at

least as high as in the solar neighborhood, i.e.,

the criteria do not have the same validity in the

galactic center as in the halo. So whatever it is we are measuring (and metal abundance is the most

likely answer) increases from halo to core as well

as from halo to solar neighborhood. Furthermore,

it begins to look as though it increases from solar neighborhood to core. I emphasize that this

is a comparative procedure and does not depend as much as Anne implies on an absolute interpretation

via model atmospheres.

Payne -Gap o s chkin : One should not define population in both ways. One must define it either by

composition or by age. A star of one-tenth the

solar mass can be 10^ years old and still be an original population star.

Smith: I think this is the source of the con87
fusion. We've become accustomed to assuming that

old stars have low metal abundances. We must now

realize that we have two independent parameters;

stars can be old and can have high metal abundances.

Payne-Gaposchkin : It might be a good idea to

scrap Populations I and II and attempt to get some- thing a little more precise. Many people base their

use of them on one criterion, implying the other,

which is extremely dangerous.

Roman : One thing has been bothering me throughout this whole discussion. Anne's argument is based

on B-stars and early A-stars and to some extent on RR Lyrae stars. The cepheid arguments, on the other

hand, are all based on later type stars. Are we not

trying to compare apples and oranges and worrying

because they don't compare very well?

Underhill : Could we possibly turn now to the

question of the evolutionary stage of Wolf-Rayet

stars? Are they in the pre-main sequence contraction phase or the post-main sequence phase? Obviously this is tied up with questions of their binary character and mass exchange , but I would also

like to consider the evolution of the single WR

star.

Payne-Gaposchkin : Let me start this discussion

by raising a question about close binaries. AO

Cass, CC Cass, and 29 Can Maj are all early type

eclipsing binaries. They all have periods of 3 or

4 days and masses of around 10 M©. The first two

spectra are 09; the third is 07f + 0. Why has not

one of these become a WR star? Is it going to become one, or has it been one in the past? What

fraction of its life should - or might - such a star spend in the Wolf-Rayet phase? From this kind

of statistics we ought to be able to verify the idea

that Wolf-Rayet stars should occur at some point in

the lifetime of early- type binaries.

Smith: For a 25 M^j primary (Kippenhahn 1968) , the fast part of the mass exchange (B to C in Figure

4) takes place in about 300 years; the slow part (C

to D) , in about 4000 years. For stars of lower

mass, it takes a little longer, but not much. Computations have been made for primaries of three

different masses: Paczynski started with a 16 M©

star and ended up with about 6 M©; Kippenhahn and

Weigert (196 7) began with 9 M® and ended with 2 Kippenhahn began with 25 M® and ended with 8.5 M®

Sahade: Paczynski and Kippenhahn started their

work on mass exchange between components of close

binaries in order to explain the Algol systems,

which are binaries consisting of a main sequence A88
type star of normal mass and a G- or K-type supergiant of abnormally low mass, sometimes as low as

0.2 M®. This is a surprising combination because,

assuming the two components of the binary were

formed at the same time, we would expect the heavier

component to have evolved faster. We conclude,

therefore, that the supergiant was initially the

heavier component but that at some stage in its

evolution it has lost mass to its companion. The

computations explain very nicely how these systems

can result from mass exchange between components of

specified initial mass and separation. Here again

the transition appears to be too fast to allow us any possibility of observing it; indeed of all the

Algol systems investigated, only one shows any sign

of being in the transition stage.

Mrs. Gaposchkin has asked about systems with

0-type components: well, if you accept a certain

interpretation of the peculiar distribution of ra- dial velocities, the components of AO Cass (09III

+ 09III) have approximate masses 18 M<5 and 16 M©;

the components of CC Cass (09IV + 09IV) have masses

19 M© and 9 M©. Could these be systems caught in

the transition phase to WR phenomenon?

Smith: Yes, that sounds possible. I would

like to say a little more about the difference between Algol-type systems and the systems we think

produce WR stars. To obtain an Algol-type system,

mass exchange must occur before the hydrogen in the core of the primary is depleted. In the Algol

stage, which occurs near the end of mass exchange,

the original primary is the less massive component,

but it still fills its critical lobe, and it still

has hydrogen core burning. This is Kippenhahn and

Weigert's case A. Now to get a WR star, we need

their case B, in which mass exchange does not begin until after the hydrogen in the core is depleted. We then get helium core burning at the

completion of mass exchange.

Sahade: On the question of mass exchange, I would like to remark that it is not just a question

of starting with two objects of different mass, com- puting what happens if mass exchange takes place,

and being satisfied if one ends up with the desired

combination of masses. As Anne Underhill remarked

earlier, you must also explain the observed spectra. I would note further that in almost all the

theoretical calculations, it has been assumed that

the mass lost by one component is acquired by the

other. That mass is in fact being lost to the system is shown by the presence of an expanding enve89
lope as indicated by, for example. He I X3888.

Again, we must take into account the velocities of

the particles which stream out of the WR star towards the companion. In comparison with those in

Algol systems , the velocities involved in WR stars

are quite large. In V44 4 Cyg , for example, they

are of the order of 700 km/sec. The third question

I should like to raise concerns the effect of ra- diation pressure from the 0-type component on the

streaming particles. HD47129, a system with an 08

com.ponent, provides evidence that radiation pressure does affect the motion of particles being lost

from the companion star. Katherine Bracher was telling me that in one of the objects she investigated, the He I X3888 absorption line, formed in

the large expanding shell, is present during only

part of the cycle. The effect of radiation pressure may, therefore, be very important.

Schild: On the subject of evolution, I would

like to discuss some recent observations of the

association Sco OBI. This association has received

little attention to date, because it is quite far

south and difficult to observe from the northern

hemisphere. It is of special interest because it

contains two WR stars and two Of stars with P Cygnitype lines. The work was undertaken jointly with

Hiltner at Yerkes Observatory and Sanduleak at Warner and Swasey Observatory. ^The investigation is

based on spectrograms at 88 A/mm taken for classification purposes at McDonald Observatory and on photometry obtained at Cerro Tololo.

The association Sco OBI lies in the tail of the

scorpion at 16^4 7^ and -41° 38'. It contains, near

its southern end, the galactic cluster NGC 6231.

The HR diagram of the association shows a turnup

at spectral type 09 - 09.5, and some 06 - 08 stars

are present. The method of spectroscopic parallaxes

gives a distance of 2 kpc and a distance modulus of

11.5 mag. From the luminosities of the early Bsupergiants, we infer that the age of the association must be about 5 million years. This agrees

well with the age given by Kippenhahn and Smith for

a model WR star that has evolved from a 25 M© star.

Figure 13 shows the HR diagram based on the new

data. From the luminosity classes, which are shown

for the supergiants, and from the absolute magnitudes, it is clear that this association is unusually

rich in luminous stars. Whereas Underhill has re- cently noted that 0-stars brighter than absolute

visual magnitude -6.0 are very rare in the Galaxy,

Sco OBI contains 5 stars of spectral type 06 - 08

90
MV

9.0

-8.0-

-7.0 -

6.0

5.0

-4.0

3.0

05 06 07 08 09 09.5 BO B0.5 8 1 81. 5 82 83

I I I i—I I I I

I \ TT"

la

•lb • lb

• lb •n

e •lb

•

t

e

•la WN.WC

I I I

Z.A.M. S.

Figure 12. The HR diagram for the association

SCO OBI shows a turn-up at spectral type 09 - 09.5,

and some 06 - 08 stars are present. The two WR

stars are located on the diagram.

brighter than -6.0. Two of these are brighter than

-7, and their hydrogen lines show P Cygni profiles;

one of them, HD152408, shows other spectral peculiarities as well.

The most luminous star in the association is

SCO. Membership of this star has been investigated by Code and Houck. On the basis of the ra- dial velocity of stellar lines as well as the intensity and radial velocity of interstellar lines,

they concluded that it is indeed a member. The

association contains three additional supergiants.

The line drawn in the region of the B- supergiants

shows where, according to the models of Stothers

,

stars of mass 30 burn helium. From this we infer that the B-supergiants in Sco OBI are of about

30 M©. Since very massive stars evolve at almost

constant bolometric luminosity, we infer that

Sco is more than 30 M©.

There is some indication that, like other young

associations, Sco OBI has seen repeated star forma91
tion events. The HR diagram shows a fairly welldefined turnup at spectral types 09 - 09.5, and a number of hotter, presumably younger stars are present. In Figure 14, the HR diagram of Sco OBI is

compared with that of the association Perseus OBI,

which contains the clusters h and x Persei. From

the comparison it appears that we can recognize in

Sco OBI a coeval group of stars, which has a turnup

at 09 - 09.5 and which is associated with the early

B-supergiants . There is also a group of bluer stars

which may define a younger coeval group.

Now Blaauw has recently suggested that if a young association contains stars of different ages,

the youngest stars will be found in the most compact

regions. If then we assume from the HR diagram that

the 06 and 07 stars are the youngest, we would ex- pect to find them in the cluster NGC 6231, the most

condensed part of the association. This does not,

however, seem to be the case. Of the 06 and 07

stars, two are found in the cluster, and the rest

are spread over the entire association. Moreover,

it does not seem possible to distinguish the 06 - 08 stars on the basis of any independent criteria

such as proper motion, radial velocity or spatial

distribution.

It is also interesting to compare the locations

of the He-burning supergiants. Observations of the

05 06 07 08 09 09.5 BO B0.5 Bl BI.5 B2 B3 B4 B5 B6 B7 B8 B9 AO Al A2 A3

I I I I II I I I I I I I I I I I I i I I I

-9.0 -

-7.0

-6.0

-5.0

-4.0

-3.0

SCO OBI BLUE BRANCH

NGC 6231 PER

h PER

J LJ LJ LJ I I I L Figure 14. A comparison of the HR diagrams

of the associations Sco OBI and Per OBI.
three clusters, x Pers , h Pers and NGC 6231 (in order

of increasing youth) , seem to confirm the prediction

of stellar models that the sequences of He-burning

stars occur at higher effective temperatures and

bolometric magnitudes the younger the cluster.

Payne-Gaposahkin : If I remember correctly,

the Scorpio association is not associated with a group of red supergiants as are h and x Persei.

Sohild: Correct. This, too ^ could probably

be predicted from stellar models. The stability

properties of these very massive He- and C-burning

stars are being investigated at Cal Tech. It appears that when stars become very much more massive

than about 15 M© , they cannot live long as red

supergiants

Payne-Gaposohkin : But are you sure the red

supergiants in h and x Persei aren't more massive

than this?

Schild: 1 believe the evolving supergiants in

h and x Persei are certainly less massive than those

in NGC 6321. I would also point out that the red

supergiants in h and x Persei are associated with

the X Persei rather than with the h Persei stellar

population. This follows from their proper motions

and from their distribution in the association; not

only do they cluster around x Persei, but two or three such stars are found in x Persei itself,

whereas none are found in h Persei.

Payne-Gaposchkin : Neither are there WR stars

or Of stars in h and x Persei.

Schild: There are several Of but no WR stars

in the Perseus association.

Payne-Gaposchkin : Are they as pronounced as the Of stars near NGC 6231?

Underhill : No, they don't look anything like

that. There are a couple in the Scorpius association with tremendous excitation, quite unlike any

other Of star I have ever seen. I have spectra of

several Of stars in the Perseus association, and

they are quite normal. But this brings us back to our original problem: what is the difference in

spectral quality between what we call a WR star,

an Of star, an O- or a B-star?

Alter: The Of stars in Scorpius certainly

don't look anything like WR stars. They have very

sharp emission lines.

Underhill : There are only a few emission lines

in these objects; not every line is in emission.

The dominant emission lines at a4485 and X4503 are still unidentified. They are sharp. WR stars have

almost the same type of spectrum, except that the

93
lines are much broader and have different shapes.

But as Sahade and I remarked earlier, the spectra

of T Tauri stars also have characteristics in com- mon with WR stars. I do not imply that T Tauri

stars are WR stars; I am just pointing out that

while WR stars have a few characteristics in common with 0-stars, they have many in common with T Tauri

stars. We are inclined to correlate WR stars with

O- and Of stars, because they are often found together in associations; but is this getting us any

closer to the physics of whatever produces the

spectrum?

Schild: I should like to go back to the actual

WR stars in this association: HD152270 is a WC7 + O

binary in the cluster NGC 6231. HD151932 is of

spectral type WN7-A, which means that it lies in

the WN sequence containing many binaries . At low

dispersion it appears spectroscopically very similar to CQ Cephei, a well-known binary. Struve did

not consider it to be a binary, although his radial

velocities do show a period of 3.3 days with small

amplitude. I think he considered the scatter in

the observations to be too large to be certain of

its binary character.

For HD152270 we can infer a reddening correction to the UBV photometry from the reddening of

nearby stars in the cluster NGC 6231. But because

of the emission lines in the WR spectrum, an additional correction must be made if we wish to find

the color of the WR continuum. Fortunately, line

emission corrections for the U, B and V bands have

been derived by Miss Pyper from Lick coude spectrograms. Since line emission corrections are not

available for HD152270, I substituted her values

for stars of the same spectral type. The final

colors I derived are:

B-V = -0.16; U-B = -1.00; My = -6 .

In exactly the same way, I have derived the

following colors for HD151932. (In fact, for WN7

stars. Miss Pyper found that it was unnecessary to

correct for emission lines.)

B-V = -0.20; U-B = -0.98; M^ = -6 . 3 . In Figure 13, the two WR stars are located on

the HR diagram according to their luminosities and

B-V colors. To relate the UBV colors to spectral

type, I assumed that the continua of these stars

have the same intrinsic colors as the continua of

94
early B-supergiants . I made no allowance for the

binary character of the stars.

Underhill : I don't see much point in locating

these systems on the HR diagram, because the HR

diagram is for single stars, and some of these may

be binaries.

Smith: Yes, but for stars whose spectra show

no evidence of a companion, the contribution of the

companion to the luminosity must be small.

Thomas: We're talking about too many things at once here: absolute magnitudes, colors, the effect

on each of binary character and how to correct for

it, and the difference between WR stars and "associated" objects with respect to all these characteristics. Could someone take these points in order

and tell us where we stand on them?

Underhill : Schild has presented information

about the Scorpius OBI association - and in particular about the cluster NGC 6231 - to demonstrate

that he knows the distance and absolute magnitudes

of some of the 0-stars to which we keep comparing

the WR stars. He used this data to point out

characteristics of high luminosity and high excitation temperature. He then considered the two WR

stars in the association and drew conclusions on their magnitude and color. The absolute magnitude

of HD151932 is -6.3, and since the companion, if

any, is very faint, the WN7 star must have an absolute magnitude of at least -6.0. The other

pair, HD152270 (WC7 + 0) , is a very different situation. The absolute magnitude for the system is

again -6.3. According to your viewpoint, you can subtract 1 mag or 0.7 mag or nothing to get an absolute magnitude for the WR star of between -6.3

and -5.3. I personally feel there is no justification for a correction of less than 0.7 mag, although

I am not sure I would go as far as the 1.0 mag urged

by Sergie.

Allev: It is generally accepted that what we

observe here are composite spectra: an 0-star plus

a companion. What we don't know is how much of the

continuum comes from the 0-star and how much from

the WR star. Just looking at them, you get an intuitive feeling that the WR star is drowned out, as Sergie mentioned, by the 0-type companion. For ex- ample, I think that the components of Y2 Velorum

cannot differ greatly in brightness because, with

the exception of the strong carbon lines, the WR

lines are swamped by the continuum. There seems

to be no general agreement on how to correct for

the continuum; if it is a substantial component

95
in the carbon stars, it may also be important in

the nitrogen stars.

Underhill : There is one way you can try to

separate the two contributions. You can take an 0-star line in the UV, where there are not too many

emission lines from the WR star, and subtract the

two spectra. You say, for example, a single 07

star would give me a line of such and such a shape;

the observed one is only half as deep; therefore I must raise the continuum by a factor of two. Jean

MacDonald Petrie tried to do this for HD193793

using Victoria spectrograms. She showed that at

longer wavelengths, the 0-star was the brighter of

the two

Stecher: ^ Puppis and y Velorum are only about

1° apart. Many people have suggested that they

might be associated and could be jointly responsible

for the illumination of the Gum nebula. C Puppis

appears just a few tenths of a magnitude fainter

than Y Velorum. The question is, what is the absolute magnitude of an 05 star? It would be interesting to know if these stars were at the same distance, and if so, why they have the same absolute magnitude.

Undevhill : I think it's fair to say that nothing much brighter than -4.5 is known. The only Of

stars for which you have definite evidence that they

are near -6.5 are those mentioned by Schild in Sco

OBI. The Of stars in IC 1805 all seem to be doubleline spectroscopic binaries. The problem is their

distance, but again they come out about -5 or -5.5.

I'm not sure what this has to do with C Puppis.

I've never seen its spectrum, so I don't know whether it is the same as an 05f star in IC 1805.

Roman: Two comments: First, the Cygnus asso- ciation - which is not so useful because it probably does extend in depth - agrees very nicely with

Schild 's results for the Sco OBI association in that

the Of stars are a magnitude or so brighter than

the WR stars. Second, C Puppis is, on classification dispersion, a perfectly normal Of star; in

fact it is one of the prototype Of stars, and it

looks very much like the Of stars in Cygnus.

Stecher: I'd like to discuss reddening cor- rections and colors. The intrinsic color of a 3x10 '*°K main sequence star is about B-V = -0.32.

Now I thought we were all agreed that WR stars are very hot. Hanbury Brown's work and our work all

suggest that the temperature of a WR star is greater than 3x10 ""^K. So I am surprised at the low B-V

colors found by Schild. Doesn't this suggest a 96
need for greater reddening corrections?

Kuhi : The WR continuous energy distribution

over a wide wavelength region is quite peculiar.

If you try to fit a blackbody temperature to any

spectral region, you find the temperature is a function of the region you choose: the longer the

wavelength at which you fit the blackbody curve,

the lower the temperature. So B-V measurements

cannot be compared with ordinary 0-stars.

Smith: My reddening corrections were determined by comparing narrow-band photometric colors

with the intrinsic colors determined for WR stars

in the Large Magellanic Cloud. I got a reddening

correction of 1.6 mag for the WN7 star and 1.1 mag

for the WC7 + 07 system. In B-V this would corre- spond to color excesses of about 0.5 and 0.4 mag

respectively.

Sohild: My color excess was 0.46 mag, so our

values are in good agreement. But there is still

a question regarding the intrinsic colors of these

stars. Are we justified in comparing the B-V and

U-B colors derived from the continua of the WR

stars with those for la supergiants? This worries

me, especially for HD151932. Figure 15 compares

the continua of two early-type stars: CQ Cephei,

which is a spectroscopic binary and has a spectrum

similar to that of HD151932, and e Orionis which is

a BO-supergiant . The stars were unreddened to the same B-V color, using Kuhi ' s photometry. The two

continua are identical over the observable spectral

range X3300 - AllOOO.

Underhill : What does that prove? For a star

as hot as e Orionis or any of the 09 supgrgiants,

the energy distribution from ly to 3000 A is very

insensitive to the model.

Schild: What model?

Underhill : The model atmosphere. Even for a blackbody, it takes an enormous change in temperature to make any significant difference in the slope

of the continuum.

Sohild: But Kuhi has shown observationally that

the WR continuum is not at all like that of either

a blackbody or a main- sequence star.

Underhill : Another objection: it has been

suggested that CQ Cephei may be two stars in a com- mon envelope. So all you've got is a body of gas

with a rather high excitation temperature and a rather high density. What is the point of comparing

it with e Orionis? Furthermore, we have no idea

how to interpret the spectra of early type supergiants or even late B-supergiants . e Orionis is

97
• HD 214419 (WN7-A) — eORI (BO la)

2.0

1.0 1.5 2.0 2.5 l/X(;3.0

Figure 15. A comparison of the continua of

CQ Cephei, HD214419 (WN7-A) and e Orionis which is

a BO supergiant.

classified as a BO , but when you come down to details, it's nothing like a BO main-sequence star.

Schild: I agree the continuum is nothing like

that of a main-sequence star; that is just the point

I wanted to stress. But I disagree that the spectrum of CQ Cephei is all that peculiar; at low dispersion it looks very much like the spectrum of

HD151932.

Underhill : But the lines on which HD151932 is

classified are sharp Of lines. They indicate a low

density, very high excitation gas. There is little

evidence that HD1519 32 is a binary; any companion

must be so faint as to have practically no effect

on what we see. I think that CQ Cephei, on the

other hand, is an eclipsing variable, one of those

"two nuclei with common envelope" efforts. You

may get low density, high excitation gas streams

around a binary system, in which case it will of course emit the same lines, but any such gas

streams have very little to do with the stars inside . 98
Kuhi : There is another interesting point about

Schild's slide. The model atmospheres with which

I compared my WR observations were necessarily the ones available, namely those for 0 and B main-sequence stars. Schild, on the other hand, is com- paring a WR star with a supergiant, and it is interesting that the continuous energy distributions

fit so well.

Kuhi and Sohild: Perhaps what we're saying is

that the WR envelope is an example of an extended

atmosphere. We are then in agreement with Anne

that in both WR stars and supergiants , we are dealing with phenomena involving an extended atmosphere.

Thomas : Can you say specifically what you mean

by "phenomena involving an extended atmosphere"? I would agree that we are likely to find excitation

levels and density scale-heights much exceeding

those which, in hydrostatic equilibrium, are asso- ciated with your continuum temperatures. Personally,

I think the whole atmospheric phenomenon will be

explained by large-scale heating and, possibly, a momentum supply. But I am concerned that when you

people speak of phenomena involving an extended

atmosphere, you are referring only to dilution effects. And while these may be present, they cannot

by themselves begin to explain the observed phenomena, particularly if you adopt the classical approach and consider only radiative excitation. But

perhaps this subject should be postponed until a later session.

Underhill: We should consider the interpretation of the UBV colors. Quite correctly we make

the best possible corrections to get intrinsic B-V

and U-V colors for the two WR stars in NGC 62 31.

Now I don't think anyone will argue strongly against

those colors; the corrections for interstellar reddening are reasonable. Then comes the question:

Do these intrinsic colors mean anything? Do they

relate to spectral types and model atmospheres? It seems to me that the conclusion reached here - which

I heartily support - is that they don't mean much.

We've got them, but we'd better not interpret them

the same way we do for 0- and B-stars.

Thomas: I'm lost: why do we need to bring in

model atmospheres at this point?

Underhill: Because the interpretation of intrinsic colors in terms of effective temperature is

done by means of model atmospheres. I'm inferring

that because Schild put his stars on an HR diagram

as one would for B-stars, he was prepared to assign

them effective temperatures. This is the normal

99
procedure when interpreting an HR diagram.

Sohild: That is indeed what I had in mind.

May I thank Anne for being so explicit.

Thomas: Then I think we have arrived at the same point we reached in our discussion of "phe- nomena involving extended atmospheres", namely that

point at which we postpone further discussion until

such time as we can be more specific about what we

are trying to do and the physical basis for doing

it.

100
PART B

A SURVEY OF SPECTROSCOPIC

FEATURES OF WOLF-RAYET STARS

CHAIRMAN: Jorge Sahade

Ohservatorio Astronomioo^ Vniversi-dad

Naoional de La Plata

INTRODUCTORY SPEAKER: Leonard V. Kuhi

Department of Astronomy ^ University

of California at Berkeley

I. Introduction 10

II. Detailed Spectroscopic Features of

Wolf-Rayet Stars lOZ

a. Identification of Spectral Lines 10 2 b. The Continuous Energy Distribution 10 8 c. Line Profiles 114

d. Line Intensities 124

e. Variations in Line Intensities

and Profiles 128

III. The Effect of Binary Nature on Spectra 129

IV. Intermediate Objects 1S7

a. Wolf-Rayet Stars Showing Lines of

Both Carbon and Nitrogen 1Z7

b. Stars Showing Wolf-Rayet Features

and Nebular Lines 158

c. Central Stars of Planetary Nebulae 141

References 142

DISCUSSION 144

101

I. INTRODUCTION

The general properties of Wolf-Rayet stars

and their spectral classification schemes have already been discussed at length by Lindsey Smith in

the first paper of this symposiiim. I would like

therefore to discuss in some detail the spectroscopic features of these stars as determined from

the observations. The outstanding characteristics

of the Wolf-Rayet stars are the extremely strong

broad emission lines of the ions of He, C, N and 0

and their dichotomy into two apparently separate

sequences: one showing predominantly lines of He,

C and 0; the other, lines of He and N. It is with

the nature of these emission lines and the occasional absorption components accompanying them that

I will concern myself. In particular their identification, variability, profiles, intensities, possible correlations, and peculiarities are of great

interest in providing clues to the understanding

of Wolf-Rayet atmospheres. In addition the peculiarities of the continuous energy distribution

will be described along with the apparent effects

of binary nature on both the emission lines and

continuum. Finally a brief description of other

objects showing some Wolf-Rayet characteristics

will be given. Most of the discussion will be

from the observational point of view, but I cannot

promise to restrain myself on the theoretical implications whenever I feel these to be of great importance.

II. DETAILED SPECTROSCOPIC FEATURES OF

WOLF-RAYET STARS

a. Identification of Spectral Lines

The stronger features of the emission-line

spectra have long been adequately identified. Wavelength lists in the normal photographic region were given originally by Deals (1930) and more re- cently by a variety of authors: Swings (1942) , Aller (1943), Smith (1955), and Underhill (1959, 1962,

1968a) . The extension of the observations into the

near-infrared by Swings and Jose (1950) revealed a wealth of new lines demanding identification, especially in the carbon (WC) sequence. Edlen (1956)

103
used new laboratory data for C II and C III, together with predicted wavelengths from hydrogenic

transitions expected from C IV, to successfully

complete the identification of nearly all the lines

of any consequence showing up in Wolf-Rayet spectra.

He also pointed out the remarkable extent to which

the spectrum of carbon (especially C IV) is developed in the WC stars. The nitrogen spectrum is

rather bleak in comparison. Additional work in the

photographic infrared (i.e., out to X8600) was re- ported in a series of papers by Andrillat (1952,

1953, 1957, 1958, 1962) along with that in

Smith's thesis (1955). Miller (1954) made use of

very low dispersion objective-prism spectra to

reach the l]i region and succeeded in finding the

extremely strong line of C III A9710 in the WC

stars. Code and Bless (1964) obtained a low dispersion spectrum of yz Vel (WC7 + 07) and identified most of the emission lines as being due to

C II, C III, He I and He II. With the advent of

photoelectric spectrum scanners with photomultipliers capable of observing out to ly with no great

difficulty, the infrared spectra of five Wolf-Rayet

stars were obtained by Kuhi (1966a) . The wave- length region from X8000 to XllOOO was scanned with

a 10 A exit slit, and typical results are shown in

Figures 1 and 2. The outstanding points are the

richness of the spectra of C ions in the WC stars

as compared to the absence of N features in the WN

stars. In contrast, the WN stars show only very

strong lines of He II (especially A10124) and He I X10830. In addition, following Edlen's work on

C IV and using the laboratory data of Glad (1953)

and Bockasten (19 55) , Kuhi also succeeded in identifying most of the weaker lines in this spectral

region. In the WN stars one line at X10430 re- mained unidentified, and N III was tentatively suggested as an identification pending further laboratory work.

The results of these investigations can be

summarized as follows. Firstly, the spectra of

Wolf-Rayet stars represent a very high degree of

excitation and ionization. Lines of He I and He II are present in both WN and WC stars, although He II seems to be relatively stronger in WN's than in

WC's. In addition, lines of H I are usually over- shadowed by those of He II (except in a very few

stars) and apparently contribute little to the

spectrum. For example the alternate members of the

Pickering series of He II blend with lines of the

Balmer series of H I. Yet in most stars (e.g..

104
8000 8500 9000 9500 10 000 10 500 II 000

X Figure 2. The infrared spectrum of HD192163 obtained as in Figure 1. The break at A10530 is due

to a change in integration time.

105
HD191765, Underhill 1968a) the Pickering decrement

is quite smooth so that any hydrogen contribution

must be negligible. Secondly, the spectra of WC

stars are dominated by the lines of C II, C III and

C IV along with weaker lines of O II, 0 III, and

O IV, The strongest features in a typical WC7 star

are C III XA4650, 5696, 9710, C IV AA4650, 5808,

7726 and He II A4686. No strong lines (or even

lines of moderate strength) of any ion of nitrogen

appear in WC stars. The spectra of the WN stars

are not quite so spectacular, having only moderately strong lines of N II , N III, N IV and N V. The

strongest features are usually He II AA3203, 4686,

8242, 10124 (A9345 is located in a region of very

heavy atmospheric water-vapor absorption, and even

though clearly present its strength is difficult

to measure precisely), N III AA4100, 4640 and N IV

A7112. Again in general there are no strong lines

of ions of carbon or oxygen present in the WN stars,

with the sole exception of C IV A5808 (see below)

Thus the two classification sequences (WN and

WC) are dramatically borne out by the entire

spectral region from A3100 to A 11000. A few south- ern stars (classified as Wolf-Rayet stars) which

seem to violate this general rule will be discussed

later.

We may now consider some finer points in more

detail and try specifically to answer the question

which has been consistently (and persistently)

raised by Underhill (196 8b) . Is there any concrete

evidence for the presence of N in WC's and C in

WN's? Let us look first at the line at 5808 A

which occurs with moderate strength in WN's and is one of the stronger lines in WC's. This line has

been variously identified as being due to C IV and

to N IV. Both Swings (1942) and Aller (1943) identified the line as a blend of C IV AA5801, 5812.

Because Hiltner and Schild (1966) have ascribed it

to N IV it is perhaps worthwhile to reproduce

Swings' original arguments and settle this issue

once and for all. The predicted C IV lines arise

from the transitions 3s ^S - 3p with wavelengths 5801.3 and 5812.0 A and intensities 10 and

8 respectively, with the excitation potential of

the upper level being 39.5 eV. The blended wave- length would be 5 806 A, in good agreement with

that observed. However the predicted N IV lines

involve the transitions 3p' - 3d' ^P° (with an excitation potential of 63.1 eV) , giving rise to

wavelengths 5812, 5828, 5846 A with intensities 5,

5, 15 respectively (according to Swings). The ex- pected blend occurs at 5 830 A which does not agree

106
with the observed wavelength nor can one expect

such a high excitation line to appear with much intensity. In addition, recent laboratory work by

Hallin (1966) indicates that this predicted multiplet of N IV would be far too weak to have any no- ticeable effect on the spectrum. Therefore we can only conclude that this line is due to C IV X5806

and hence that the C IV ion does appear in WN stars

(see also Underhill 1968a)

However when we consider the evidence for

other ions of carbon or oxygen in WN ' s and of nitrogen in WC's, the situation is no longer so unam- biguous because we are forced to deal with very

weak lines or blends with much stronger lines. For

example in her paper on HD192103 and HD192163

Underhill (1959) states that N III is definitely

present in the WC7 star. However the evidence

comes entirely from such statements as "rather too

strong for 0 III, C II, C III, and C IV". There

is no clearcut case of an N III line free from

blends. She also states that in the WN6 star, 0 III

is definitely present and O IV is probably there.

Here the evidence is somewhat more convincing,

namely weak lines at A5506 and X4797, but it is

still not conclusive. Andrillat (1952) has also

managed to identify N II in WC stars and C I through IV in WN's. As Edlen points out, the

strongest lines of N II in the infrared are XX1162

and 8439. These lines have not been observed in

normal Wolf-Rayet stars, neither in WN's nor in

WC's. Furthermore the identification of C III in

WN's is extremely tenuous. The far-infrared work

has clearly shown that the dominating feature in

WC's is C III A9710 (often with a central intensity

greater than ten times the adjacent continuum) and

that there is no trace or hint of such a line in

the WN's. Therefore we can say definitely that no

C III occurs in WN stars. Underhill (1968a) concurs with this and in fact leaves a weak feature at

A5700 in HD191765 unidentified because of the ab- sence of X9710 in the infrared.

We are then fairly safe in concluding that for

the "normal" Wolf-Rayet stars the dichotomy is complete. (The somewhat "peculiar" Wolf-Rayet objects

will be discussed later.) These results imply a strong argument in favor of abundance differences

being responsible for the two vastly different spectral sequences. Underhill (1957) has, however, suggested that the differences can be interpreted in

terms of differing excitation conditions, brought

about by a completely different far-ultraviolet en- ergy distribution for the WN's and WC's. One point

107
is still extremely difficult to explain with this

picture: How is it possible to construct an atmosphere which is able to distinguish between the

47.24 eV ionization potential of N III and the

47.67 eV of C III so completely that no lines of

N III are seen in WC's and no C III in WN's? It is

inconceivable to me that no trace of C III X9710

would be found in WN's if the separation into two

sequences is really due to a difference in excitation conditions, since the differences between C III

and N III are so minute. The only plausible ex- planation so far offered is a bona fide abundance

difference of carbon and nitrogen. An atmosphere

designed to distinguish 47.24 eV from 47.67 eV

seems very artificial.

The wavelength identification problem seems now to be basically solved. However a niimber of

weak lines still remain to be identified: examples

are XX7426, 6503, 5700, and weaker lines in the

ultraviolet which occur in WC stars, and AA10430,

5200 in WN's. Presumably further laboratory work

on the spectra of multiply ionized ions of common

elements will solve this minor problem in the near

future

h. The Continuous Energy Distribution

I would like to discuss here only the recent

photoelectric measurements of the continuous energy

distribution and to refer you to the literature

(cf . , the work by Andrillat and Vorontsov-Velyaminov

cited by Kuhi 1966a) for earlier photographic results. The outstanding feature of the photographic

work was the extremely low values of color temperatures obtained for the Wolf-Rayet stars, i.e.,

Tc 7x10^ to 1.6xio'*°K. These temperatures seemed

to be incompatible with the high excitation and

ionization temperatures indicated by the emission

lines. In order to clarify the situation I undertook a study of the brighter Wolf-Rayet stars ac- cessible to Mt. Wilson, using the photoelectric spectrum scanner on the 60- and 100-inch telescopes.

The work has been described in detail (Kuhi 1966a)

in the literature so that we need only outline it

here. Exit slits of 40-50 A were used in the blue

and 50-70 X in the infrared. Great care was taken

in choosing the continuum wavelengths to avoid the

emission lines , but even so it was found necessary

to make some corrections for their presence in the

WC stars. Reddening corrections were made by means

of a narrow-band three-color system; i.e., the early

type stars lay on a vertical line on a color-color

108
plot and the Wolf-Rayet stars on two essentially

parallel reddening lines, the WC's being cooler than

the WN's. It was assumed then that the Wolf-Rayet

stars could be unreddened to the same point as the

location of the early type stars in the two-color

diagram. The results can be summarized as follows:

1. The energy distributions of Wolf-Rayet

stars do not resemble those of ordinary

O- and B-stars.

2. There is a slight ultraviolet excess

for the WN's.

3. There is a strong infrared excess for

both groups, but especially for the WN's.

4. The WC's seem to be somewhat cooler than

the WN's.

The net result of this peculiar energy distribution is that the color temperature assigned to the

star depends on the wavelength at which it is measured: The longer the wavelength the lower the

color temperature. Typical values for a WN star

range from >10^°K at A3500 to l.lxlo'*°K at X9500,

and for a WC star, from 6.5xl0'*°K at X3500 to

1.5x10** °K at X9500. Consequently it proved impossible to fit any theoretical energy distributions to

the observations. The stars are clearly quite peculiar. This peculiarity is also evident when one looks at the energy distributions for binary stars

in which the OB-type comparison contributes a large

fraction of the energy observed. For example

HD193514 is dominated by the 0-star and hence the

continuous energy distribution looks like that of

an 0-star, whereas in HD214419 both members are equally bright and the resultant distribution looks

much more like that of a single Wolf-Rayet star,

i.e., quite steep in the ultraviolet and quite a bit flatter in the infrared than for normal 0-stars.

Since the 1966 work was published, Hayes (1967)

has recalibrated the absolute energy distribution of

Vega and other standards. He finds a somewhat

larger Balmer discontinuity and a steeper slope in

the infrared than was shown by the Oke (196 4) calibration on which the Wolf-Rayet continuous energy

distributions were based. Since this new calibration has led to much more consistent results for Aand B-stars, I have applied the correction factors

to the 1966 data and present the results in Tables

1 and 2. These tables give the unreddened fluxes

in mag per unit frequency interval. A few typical

curves for single and binary stars are plotted in

Figure 3 along with an 09V and a BOIa star. The new calibration has two effects on the earlier re109
TABLE 1

UNREDDENED FLUXES: WN STARS

186943 190918

HD4004 9974 50896 WN4 187282 WN4 .

1/X WN5 WN3 WN5 +B WN4 +09.5Ia

3.012 + 0.18 + 0.46 .+ 0.47 + 0.52 + 0.18 + 0.84

2.932 0. 24 0.49 0.46 0.59 0.24 0.55

2. 800 0.35 0.56 0.57 0.62 0.41 0.58

2.750 0.38 0.65 0.62 0.67 0.44 0.67

2.632 0.40 0.58 0.56 0.58 0.45 0.57

2.545 0. 49 0.70 0.61 0.63 0.53 0.60

2.522 0.50 0.72 0.59 0.65 0.53 1.64

2.410 0.66 0.75 0.69 0.73 0.66 1.65

2.350 0. 75 0.84 0.75 0.76 0.72 1.68

2.259 0.82 0. 84 0. 80 0.83 0.78 1.80

2. 089 0.91 1.01 0.92 0.94 0.85 0.94

2.000 1.00 1.00 1.00 1.00 1.00 1.00

1.949 1.01 1.12 0.96 1.02 1.07 1. 05

1. 900 1.11 1.06 1.05 1.07 1.15 1.04

1. 866 1.35 1.16 1.11 1.17 1.16 1.16

1. 800 + 1. 36 1.23 1.11 1.18 1.17 1.17

1.681 — 1.32 1.22 1.31 1.32 1.32

1.570 1. 44 1. 30 1.42 1.47 1.40

1.471 1.57 1.37 1.53 1.54 1.62

1. 322 1. 64 1.51 1.52 1. 87 1.76

1.190 1. 87 1.62 1.83 1.99 1.93

1.049 2. 06 1. 76 1.98 2.19 2.14

1.012 2.13 1.78 2 .17 2.25 2.19

0.934 + 2.09 1.85 + 2.11 +2.38 2.33

0.900 + 1.90 + 2.44

Note: The unreddened flux is given as -2.5 log

Fv + const, and is normalized to 1.00 at

A5000.

110
TABLE 1 (cont.)

UNREDDENED FLUXES: WN STARS

193077 193576 193928 211853 214419 219460

191765 1917 65 WN5 WN5 WN5 WN6 WN7 WN4 .

WN6 WN6 (+0B) +0 6 +0B +BOI +07 +B0

+ 0 . 3o + 0 . 41 _i_ A + U . A C 45

0.38 A IT A 0.50 A0 . C A 5 9 A C Q U . 5 o

AU . a A o U 0 . 51 A C A AU . a c. DO

r\ CO U . OZ A C /I A c. a D o

U . D U

A an U . D / A D 1

U , D J

A A A

0 . C A D U U . /D A 1 "i A

U . D / A "7 0 0,72 A O T 0 . o7 A0 . 6 9 n o o A Q 1 u . y 1

AU . ol

1.00 1.06 0

92

1.00 T A A 1.00 1 . A A 00

T A /I 1.04 n AO 1.03 1 . A \Z Ob

1 , 11 1 . 11 1 . A A 02

1.21 1.28 1 . 11

1.27 T 1 A 1 . 30 1 . 19

1.37 1 c 46 1 . 33

1.43 1.58 1. 45

1.59 1.69 1. 61

1.74 1.88 1. 78

1.92 2.04 1. 94

2.07 2.17 2. 24

2.11 2.35 2. 26

2.13 2.45 2. 35 + 2.26 + 2.49 + 2. 38

45 +0.50

0.56

0.70

0.69

0.63

0.61

0.63

0.70

0.71

0.79

0. 95

1.00

1.07

1.09

1.17

1.26

1.35

1.45

1.51

1. 80

1.94

2.08

2.25

2.43

38 +2.54

+ 0.48

0.54

0.71

0.61

0.64

0.68

0.70

0.66

0.75

0.76

0.94

1.00

1.02

1.06

1.15

1.17

1.28

1.38

1.48

1.64

1.79

2.19

2.13

2.32

+ 2.29

+ 0.43

0.63

0.62

0.65

0.63

0.66

0.72

0.71

0.74

0.82

0.87

1.00

1.06

1.08

1.15

+ 0.43

0.44

18

30

40

56

69

84

46

51

52

58

61

65

68

78

93

00

99

2.20

2.18

2.32

+2.38

1.10

1.19

1.13

1.35

1.44

1.61

1.86

1.93

2.11

2.26

2.43

+ 2.31

+ 0.60

0.65

0.79

0.73

0.70

0.71

0.71

0.74

0.78

0.83

0.94

1.00

0.98

1.11

1.16

1.21

1.37

1.45

1.58

1.78

1.93

2.17

2.20

2.37

+2.30

111
suits: (1) Comparisons made on a relative basis are not changed, i.e., the observed energy distribution

of a single Wolf-Rayet star is still peculiar when

compared to the observed distribution of a normal

0-star. The ultraviolet and infrared excesses clearly remain. (2) Calculations and comparisons made on

an absolute basis are quite different, i.e., the es- timates of color temperatures will all be modified;

being slightly lower in the ultraviolet and slightly

higher in the infrared. The overall effect of the

dependence of color temperature on wavelength is

still present. The disagreement between the observations and the predicted energy distribution for

Tgff % 5xl0'*°K and log g = 4.0 is now somewhat

smaller in the ultraviolet and photographic but is

TABLE 2 UNREDDENED FLUXES: WC STARS

168206 19 2641 19 3793

HD16523 17638 165763 WC8 19210 3 WC7 WC7

1/ A +B0 O WCo +Be +05

3.012 + 0.25 + 0.35 + 0.47 +0.33 + 0.45 0.48 0.52

2. 850 0.55 0.66 0.60 0.40 0.73 0.65 0.83

2.734 0.51 0.51 0.53 0.51 0.64 0.62 0.67

2.621 0.54 0.62 0.60 0.53 0.54 0.54 0.57

2.484 0.62 0.65 0.70 0.62 0.57 0.62 0.64

2.342 0.67 0.66 0.67 0.69 0.68 0.74 0.75

3.025 0.97 1.02 0.96 0.97 1.00 0.90 0.92

1.980 1.00 1.00 1.00 1.00 1.00 1.00 1.00

1.925 1.05 1.10 1.02 0.99 1.04 0.97 0.94

1. 866 1.01 1.12 1.04 1.10 1.08 1.09 0.99

1.808 1.20 1.27 1.13 1.15 1.24 1.19 1.13

1.669 1.30 1.39 1.27 1.32 1.28 1.27 1.27

1.385 1.62 1.60 1.48 1.56 1.59 1.54 1.47

1.250 1.76 1.70 1.56 1.57 1.73 1.66 1.60

1.098 1.86 1.95 1.78 1.92 1.93 1.91 1.78

1.056 2.01 2.02 1.79 1.93 2.00 1.94 1.84

0.950 +2.36 2.27 1.95 2.12 2.16 2.16 2.04

0.898 +2.26 +2.07 +2.14 +2.36 +2.08 +2.13

112
still quite large at ly ('^^30%) . Thus one can safely

conclude that the Wolf-Rayet stars are peculiar in

their continuous energy distributions when compared

to normal OB main sequence stars. However as shown

in Figure 3, the differences between Wolf-Rayet

stars and OB supergiants may not be so great. This

point is in need of further investigation.

Finally we note that the depth of secondary

eclipse in V444 Cygni (0-star in front) in the

continuum is an increasing function of wavelength

ranging from ^ 0.14 mag at A3300 to ^ 0.19 mag at

XllOOO. Kuhi (1968) has shown that this result is

most readily accounted for by assuming that the

Wolf-Rayet star is intrinsically brighter in the

infrared than the 0-star. He is able to reproduce

reasonably well the dependence of eclipse depth on wavelength by using the observed energy distribution of a single WN5 star to compute the brightness

0.5

1.0

1.5

2.0

1 1 1 ^ 09 V-- ^XBOIo

NX \ X*

\\ \ v\ \ \ \v \ \ \\ \ \ \\ \ \ \ \ \ \ \ K \ \ \ _ \ ^\ \ \ \

\ \ \

1 1 1 ^ i\

3.0 Figure 2. The continuous energy distribution of

a WN6 star compared with that of the 09V star 10

Lacertae (thin dashed line) , the BOIa star, e Orionis (thin solid line) , and an unblanketed model

atmosphere for T^ff = 5x10 '*°K (unlabelled thin

solid line) . The units are mag per unit frequency

interval on the ordinate and inverse microns on the

abscissa.

113
ratio as a function of wavelength. Therefore we

are able to conclude that the infrared excess is an intrinisic property of Wolf-Rayet stars.

o , Line Profiles

The emission-line profiles for both WC and WN

stars fall into two groups, those that have violetdisplaced absorption components and those that do

not. The lines of the first group seem to show

some tendency toward the classical flat-topped profiles expected from a spherically symmetric envelope expanding with constant velocity (see Beals

1930) , but in general the lines of both groups show

rather steep sides with rounded tops. The determination of these emission-line profiles by ordinary photographic means is extremely difficult for

the stronger lines. The intensity ranges encoun- tered (e.g. , central intensities 10 or more times

greater than the adjacent continuum) greatly sur- pass the workable range of photographic emulsions,

and hence most of the photographically determined

profiles cannot be very accurate. No such problem

presents itself for the lines of moderate strength,

and photographic profiles should be quite reliable.

For the stronger lines photoelectric spectrum

scanning provides a convenient and reliable means

of obtaining profiles. Unfortunately very little

work has yet been done in this area: my own results

for a few lines seem to be all that is available

with sufficient resolution to be at all comparable

to the photographic profiles. Figures 4-6 illustrate photoelectric emission profiles for the

X4600-X4700 region for a number of Wolf-Rayet stars.

They were obtained with a 2 A exit slit with the

Oke scanner used on the 100-inch reflector at Mt.

Wilson. The extreme range in intensity is clearly

evident: the peak intensity of C III A4650 in

HD165763 is 20 times the continuum intensity; the

total width at half-intensity is 55 A. The spectral types on the figures are on the Beals (19 30)

system. It would be more appropriate to use the newer system of Hiltner and Schild (1966) or L.

Smith (1966) : The three WC stars would then be

HD165763 (WC6), HD192641 (WC7), and HD192103 (WC8 pec) on the former system, and WC5, WC7 + Be, and WC8

(+ OB) on the latter. In either case a clear trend

in increasing width of the A4650 feature with earlier type is apparent from Figure 4. There is also

evidence for a displaced absorption component in

114
HD192103 even at this rather low resolution (2 A)

Figure 5 shows three WN stars HD193077 (WN5 + OB)

,

HD192163 (WN6) and HD191765 (WN6) (L. Smith's spectral types) . The important things to note are the

changes in the relative strength of N V X4605-22

and N III X4640 between types WN5 and WN6 and the

lack of any apparent correlation between the line

width and spectral type. Figure 6 shows three more

stars, HD193793 (WC7p + 05), HD190918 {WN4 . 5 + 09.5

la) and HD228766 (WN7 + O) (L. Smith's spectral

40000

20000

30 000

20000

10000

1 -

1 1 1

- / A V

' 1 \. HDI65763

\ WC6 -

\ \ HD 192641

/ \ WC6

[/ \ HDI92I03

\ WC7

1 1 1 1 4600 4700

X

15000

10000

- 5000

-0

Figure 4. Photoelectric line profiles of X4650

complex for WC^^stars obtained with a spectrum scan- ner using a 2 A exit slit. The ordinate gives the

total number of counts (proportional to intensity)

the scales on the left refer to HD165763 and

HD192103 and that on the right refers to HD192641.

115
U W u u

1

1 1 1

1 0 000 ^^-^^^^^ 91 765

WN6 1 uu u

10000 - J\ - / ^^v_y \ HDI92I63

10000

5 000 0

0 WN5

1 1 1 1 4600 ^ 4700

Figure 5. Photoelectric line profiles of A4650

complex for WN stars. Other data same as Figure 4.

types) . Again evidence for an absorption component

is visible in HD193793, and the profile itself seems almost flat-topped* For the other two stars there

is clearly no advantage in using photoelectric

methods since considerable detail is lost because

of the low resolution.

Finally we present photoelectric profiles for

the He I A10830 region obtained with a 4 A exit

slit using the prime-focus scanner on the 200-inch

reflector at Palomar. The stars scanned were HD192103 (WC8 + OB) , HD192641 (WC7 + Be) , HD192163

(WN6) and HD193077 (WN5 + OB) ; the results are shown in Figures 7 and 8. All the lines show the

strong violet-displaced absorption core similar to

that seen in He I X3889. But in addition we note

that the four stars have four completely different

profiles: HD192163 has the classical flat-topped

profile, HD192103 has a profile similar to non- metastable lines, and the other two show something

in between. There also seems to be some central

fine structure which is greater than the expected

fluctuations from the counting statistics. In this

116
10 000

5000

0 —

HDI909I8

WN5 + 09.5

10000

5000

5000- 0

HD228766

WN7

0

4600 4700

X Figure 6, Photoelectric line profiles of X4650

complex for other WR stars. Other data same as Figure 4

region of the spectriim the advantages of the scanner

technique are obvious, and it is to be hoped that more accurate profiles of the strong infrared lines

will soon become available. Good examples of photographic profiles have been given by Underhill (1968a,

1959) and are illustrated in the figures accompanying her presentation at this symposium. References

to earlier work are given in Underhill 's review paper (1968b).

One can summarize the photographic and photoelectric data by saying that the typical emissionline profile is quite symmetrical and is of the

form exp(-AX^/AXo) / where AAq is a suitable con- stant for the line and is related to the velocity

of the atoms. In general for lines of the same

ion, the width at half-intensity is proportional to

the wavelength, e.g.. He II XX3202, 4686 and 6560

in HD191765 are 25, 41 and 55 A wide respectively.

This implies that AX/X is approximately constant

and hence that the chief broadening mechanism is

simply Doppler broadening due to the motions of the

117
2000-

lOO'O

HD 192103

WC7

HD 192641

WC6

4000

3000 u

2000 (ou

1000

10700 10800 10900

X

1 1000 Figure 7. Photoelectric line profiles of X10830

for WC stars obtained with a spectriim scanner using

a 4 A exit slit. The ordinate gives the number of

counts per second.

atoms. The corresponding velocities for these

three lines are 2340, 2620 and 2590 km/sec res- pectively, which indicates that the correlation is

not one to one and hence that some other broadening

agent is also operating. Beals pointed this out in

1929 when he noted that the width increases with

wavelength at a rate slightly greater than predicted by pure Doppler broadening. In addition the

wings of lines such as He II X4686 and X3203 in

HD192103 and HD192163 (Underhill 1959) are somewhat

greater than those of a Gaussian profile, again

suggesting the presence of an additional broadening

agent. Stark broadening has been proposed (Johnson

1954) but is clearly ruled out for several reasons:

(1) Lines of different series of He II would have

different shapes and widths, since the Stark effect

would be greatest on those lines having a larger

value of n for the lower state. However no such

118
2000 0

in

in

HD 193077

WN5

<^ 1000

° 1 0700 10800 10900 1 1000

X Figure 8, Photoelectric line profiles of X10830

for WN stars. Other data same as Figure 7. effect is observed. (2) Hydrogen-like transitions

(i.e. , between principal quantum numbers) are observed to quite high values of n in C III and es- pecially in C IV, which has an unusually well-developed spectrum. For example C IV lines from the

6-n, 7-n, and 8-n series are identified in WC stars.

In the presence of any appreciable ionic field these

lines would not be visible since they are exceedingly sensitive to Stark effect.

The work of Miinch on the absorption lines of

the 0-type companion of V444 Cyg (1950) gives con- vincing evidence of the presence of electron scattering in Wolf-Rayet atmospheres. This had already

been suggested (Kopal and Shapley 194 6) from the

eclipse data obtained by Kron and Gordon (1950)

,

which revealed a very broad primary eclipse com- pared to a much narrower secondary eclipse. Munch

noted that the absorption lines from the 0-star

were much broader and shallower when the Wolf-Rayet

star was in front (primary minimum) than at other

phases. He interpreted this as being due to electron scattering both in the Wolf-Rayet envelope and

119
in the 0-star atmosphere itself. An optical depth

T - 0.5 for the scattering "layer" was found nec- essary to match the observed changes in the profile

of HIO. Munch also suggested that electron scattering produces an appreciable broadening in the

emission lines formed in the deeper layers of the

Wolf-Rayet envelope where the optical depth would

be sufficiently large. In this way he is able to

reproduce roughly the profile of He II X4686 in

V444 Cygni with electron scattering producing about

one-half of the observed width.

Underhill has criticized this result on sev- eral occasions (1959, 1968a) by stating that any

absorption line should undergo similar broadening.

For example the violet-displaced absorption com- ponents of He I (and for that matter any of the

observed absorption components) are much sharper

than any emission line in the same star. The most

likely conclusion is then that the absorption com- ponent is formed in a part of the envelope that is

basically outside the electron scattering region.

This she dismisses as an unlikely hypothesis. How- ever I would like to suggest that this is not so unlikely a possibility, especially for He I XX3889

and 10830. Firstly, these absorption lines arise

from metastable levels 2 - 3 and 2 - 2 respectively; as is well known, these lines (as well

as others of higher excitation) are considerably

enhanced under conditions of moderate dilution

(W <. 0.01) because of the buildup of the population

of the 2 metastable level. Secondly, in the

region where electron scattering is efficient (i.e.,

T - 0.5) , the material density must be fairly high

10 ^^/cm^) so that collisions must play a significant role (the importance of collisions has been

discussed most recently by Code and Bless 1964 and

Bappu 1967) and hence act to depopulate the metastable level. These two factors imply that the

He I absorption lines are formed at some distance

from the star and quite possibly outside the bulk

of the electron- scattering region. The other absorption components that are most often seen ar^

additional lines from the He I triplets and the

blends of C III X4650 and N IV X3483. These all

arise from levels which can be expected to be over- populated by dilution effects, and again the absorption features are quite sharp in contrast to the

broad emission. Sharp absorption components are also seen occasionally for N V X4609 and C IV

X5806; these lines arise from normal levels of

quite high excitation, and Underhill (1968b) suggests that they are due to large gf values and the

120
large abundance of the ion. The latter arises be- cause of the high ionization potential required to

reach the next stage of ionization. Flat-topped

profiles, when observed, are usually found among

these lines. However the four different He I A10830 profiles illustrated earlier indicate that

we are not dealing with the case of a transparent

envelope expanding with constant velocity. The

presence of the "sharp" absorption components also

implies that the extremely chaotic conditions closer

to the star have been somehow ironed out and that

we are left with a net outward flow of material

moving at velocities ranging from 500 to 2700

km/sec. These velocities are similar to the large

velocities of expansion ('^^ 2000 km/sec) derived

by Morton (1967) from the violet-displaced absorption components observed in the resonance lines of

C IV, N V and Si IV in OB supergiants . Perhaps the same basic mechanism is at work here, but the con- ditions in the Wolf-Rayet inner envelope are clearly much more extreme.

We should now discuss the various correlations

found for the emission-line widths. For a given

star, lines from highly ionized ions are much nar- rower than those from ions of a lower degree of

ionization. In WN stars for example this is very

clearly evident when one compares N V A4609 to

N IV X3483. The correlation is not so well defined

for lines of different elements. This range in

width among different ions of the same element is

most readily interpreted in terms of stratification

in the Wolf-Rayet atmosphere. However the nature

of this stratification has not been settled. Beals

originally suggested that the temperature decreased

outwards, so that the N V lines (the narrowest) were formed in -much deeper layers than those of

N III (quite broad) . This in turn implied that the

atoms were accelerated to higher velocities at

greater distances from the star. Miinch's discussion of electron scattering however implies just

the opposite; a large optical depth produces a large

width, hence the broadest lines must be formed close

to the star, and the temperature increases outwards.

Since electron scattering does not account for the

entire width (the remainder being Doppler) , this

picture also implies decelerating motions in the

envelope. Thomas (1949) has proposed an envelope

supported by large-scale turbulence, in which the

lines of highly ionized ions occur higher in the

atmosphere than those of less highly ionized ions.

A more attractive picture (because it does not predict unobserved occultation effects) has been

121
suggested by Code and Bless (1964) . They propose

that the line widths are due to the intrinsic ve- locity range in large prominences and that collisional ionization and excitation are extremely important. The most highly ionized ions then would

have the lowest kinetic energy and hence the nar- rowest emission lines. No stratification is needed. These models will be discussed in detail by

Underhill in the following paper.

The second correlation is well defined only

for the WC stars: the earlier the spectral type the

broader are the emission lines. For example the

widths of C III-IV A4650 are 85, 45, 35 and 10 A

for L. Smith's types WC5 , WC6, WC7 and WC9 re- pectively. Presumably this correlation is re- lated to the increase in temperatures encountered

with earlier spectral types. However what the

exact mechanism or connection may be is not at all

clear. It cannot be due to higher temperature:

the widths are much too large for this if the temperatures involved are of the order of 5x10** to

10^ °K. The higher excitation (earlier type) and

larger widths may be caused by the same primary

mechanism; for example increased fliix of high en- ergy particles, as already suggested by Code and

Bless (1964) , could impart high kinetic energies

(hence Doppler motions) to all stages of ionization while also raising the general level of ionization. Clearly a great deal of work remains to

be done in this area. There is no clear cut correlation for WN stars although L. Smith (1966) does

note that "the strength and width of lines vary

markedly from class to class, increasing from WN7

through WN8 and WN6 , reaching a peak between WN6

and WN5 , and declining through WN4 and WN3". How- ever the variation from star to star and from ion

to ion is large so that no definite conclusions can

be reached.

The third correlation was discovered by Hiltner

and Schild (1966) in setting up their new classification scheme. The WN stars fall into two groups,

one with broad emission lines, the other with relatively narrow lines. Aside from one star, all the

stars in the latter group are known binaries; the reverse is true in the former group. No good ex- planation is available for this correlation, although Underhill (196 8b) has suggested that the

presence of a binary companion somehow damps out

the large velocities encountered in the Wolf-Rayet

atmosphere of a single star. However this cannot

be the entire answer because of two difficulties:

(1) There seems to be no such distinction between

122
HD 9974

HD21I564 MM

Figure 9. A comparison of spectra of binary and

single WR stars taken from Hiltner and Schild (1966)

the binary and single stars of the WC sequence ; if. a binary damping mechanism were operating, surely

it would have some effect on the WC binaries as well. (2) We might expect some additional correlation between line width and orbital separation (i.e.

period) if the effect is really due to the companion. The greater the separation, the less would be

the influence of the companion and hence the broader the emission lines produced. Figure 9 seems to

show just the opposite effect, if any.

For Hiltner and Schild 's WN5-WN5.5 stars,

HD186943 shows a much broader He II A4686 line than

HD190918, even though the periods are 9.55 and 85.0

days respectively. Again in their group WN6-WN6.5,

HD193576 (P = 4.21 days) shows a somewhat broader

He II line than HD211853 (P = 6.69 days). Finally

in group WN7, CQ Cep (P = 1.64 days) shows a broader line of He II than HD197406 (P = 4.32 days).

Thus a reverse correlation can be established ; namely, the shorter the period of a binary in a given spectral class, the broader its emission lines

123
There are therefore two counteracting effects: (1)

The narrower lines displayed by WN binaries in

general and when compared to single 'WN's, and (2)

the broader lines displayed by the close binaries

in a given spectral class. Interaction effects

obviously become more important for close binaries,

but no real understanding of their result on WolfRayet stars exists.

d. L-ine Intensities

The test of any theory for the Wolf-Rayet

atmosphere is the reproduction of its emission-line

spectrum, i.e. , not only the line profiles but also

the relative intensities of the lines from different

ions. Therefore it is of great importance to know

the observed emission-line intensities with good

accuracy. Unfortunately the problems of large intensity ranges, already discussed in connection with

line profiles, also arise in the photographic mea- surements "of line intensities. There is in addition

a severe blending problem (especially for the WC

stars) due to the presence of many weak lines, so that it is extremely difficult to isolate the intensities of individual lines. The large range in

intensity can be handled photoelectrically (although very little has been done) , but there is no easy solution for the blending. Consequently the

meaning of line-intensity lists is not too clear.

Nevertheless comprehensive studies have been

carried out by a number of people (Plaskett 1924,

Deals 1930, Swings 1942). Aller (1943) gives

photographic intensities for eight Wolf-Rayet stars

for the spectral range X3200 to X6100. Swings and

Jose (1950) extended the work into the photographic

infrared ('\^ X8600) for eight stars (but only

three in common with Aller) . H. Smith (1955) presented a great deal of data for southern Wolf-Rayet

stars to supplement that from the fairly well studied northern stars. He also pointed out the notoriously poor agreement among line intensities determined by three different observers for the same

star, factors of two being not at all uncommon. More recently Underhill (196 8a) has provided line

intensities for HD191765, and Kuhi (1966a) has determined photoelectric intensities for a few stars

in the far-infrared (X8000 to XllOOO) . Most of the

other work has been qualitative with the strengths

usually listed as weak, moderate, strong, etc. We

list in Tables 3 and 4 some new photoelectric measurements of the intensities of the stronger far124
TABLE 3 INFRARED LINE INTENSITIES: WN STARS

Spectral He II He II He I Py

Star tvDe 8240 10124 10830 10938

HD 4004 WN5 58 384 384 93

6327 W(He) 19 225 42

9974 WN3 6 116

50896 WN5 60 371 294

56925 WN5 58 346 106

65865 WN4 .

18 84 33

165688 WN6 40 317 403 170

MR 89 WN7 26 106 400 46

HD 177230 WN8 19 33 384 36

186943 WN4 + B 17 110 47 22

187282 WN4 179 38 23

190918 WN4.5 + 09. 511 38 15 23

191765 WN6 47 322 297 64

192163 WN6 45 285 326 69

193077 WN5 (+0B) 13 97 66 17

228766 WN7 + 0 1 29 41 14

193576 WN5 + 06 14 150 882 29

193928 WN5 + OB 14 221 161 32

197406 WN7 12 38 69 9 MR 114 WN5 + OB 27 156 81 32

HD 211564 WN3 25 180 19 8 211853 WN6 + BOI 11 91 91 16

MR 119 WN8 9 8 166 47

HD 214419 WN7 + 07 8 99 56 8 219460 WN4.5 + BO 11 64 32 8 Notes: 1. The spectral types are those of Lindsey

Smith (1966) and are somewhat different from the

older lAU types.

2. The intensity is given in terms of

the continuum flux per unit wavelength at the central wavelength of the line and refers to the total

intensity above the continuum measured in the exit

slit (96 A for A<9344 and 128 A for X>9344) . Thus

it may include a contribution from nearby weaker lines.

125
M ^ I—I prj 00

M ^

H CNtH

0) O

>

H O

H CO

> o

00

^ 00

f3

Jh Q)

CO

o o CN o CM LO ro 00 CO rH ro CM

LT) <y\ LO fO LO CM lO o ro rH rH

If) CN H 00 o H

CN tH rH rH CM CM rH CM

U~) rH C?i LO 00 LO ^ rH

IT) LO CNJ rH lO fO CO CvJ CO

o 00 o CO ro

CN <^ ro

iH rH H H H

ro o 1 CM LO 1 LO CO CN] ro CO ro CM ro

I

O H CO o LO 00 o

00 CM CO H LO 00 CM CM ro o ro

LO LO H LO CN LO rH CO CM ro

00 00 00 rH H ro CO o o o

LO CN CN LO 1—

<—\ H

rH LO ro CN o ro o CO CNJ O LO CO

LO CN rsi CM CM iH ro LO

CN 00 CO rH o LO LO 00 rH CTi LO CO LO CM H ro

rH o <^ H LO ro H

iH LO CO o ro CM 00 H rH

LT) CO CN LO lO CN LO 00 o

o^ 00 CO 00 CN O LO ro CN LO o H H

' ^ o CQ Q) LO O

O PQ o +

+ + + +

<D LO CO LO r- 00 LO LO U L) U U U u u U U U U u U

12 12 IS [S [2 12 12 [2 [2 12 12 2

CO o CM ro H ro a^

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CN ro CN o o rH rH LO LO 00 CN ro LO ro

CM vo U3 o <T\ a\ H

rH rH 00 rH rH H H rH rH rH H CM Q 0^ Q Q

126
infrared lines for 38 stars accessible to the 120-

inch reflector at Lick Observatory. These intensities were obtained with the Wampl^r scanner using

exit slits of 96 A (for X < 9344 A) and 128 A (for

A > 9344 A) and represent the total intensity of

line emission in the wavelength interval measured.

The tables give this intensity in terms of the

continuum intensity (per unit wavelength) at the

central wavelength. This procedure avoids calibration and reddening problems so that the absolute

fluxes can always be obtained once the flux in the

continuum is adequately known. For example Underhill (1968a) has used Kuhi ' s (1966b) continuous en- ergy distribution to obtain absolute line fluxes

for HD191765.

We can now make use of these measurements to

estimate the total energy involved in the emission

lines as compared to that in the continuum. If we

ignore the known and suspected binaries and allow

10% correction for fainter lines not measured, we

arrive at the data in Table 5. The quantity tabulated is the ratio, r, of line flux to continuum

flux in the region A 8000-XllOOO . The ratios for

each spectral class are the means for all the single

stars measured in that class. A number of hitherto

single stars were not included in the means because

their intensities seemed too low; they may be binaries even though they have not been detected as such

in the normal photographic region of the spectrum.

TABLE 5 FRACTION OF ENERGY IN EMISSION LINES

(A8000-A11000)

Spectral

type

Niamber

of stars r WN3 1 0.088

WN4 1 0.095

WN5 3 0.261

WN6 3 0.291

WN7 1 0.212

WN8 1 0.173

WC5 3 0.519

WC6 2 0.499

WC7 1 0.351

WC8 1 0.486

WC9 2 0.298

127
Examples are HD9974, 65865, 197406 and MR119. It

should also be noted that two o£ the WC5 stars have

20% lower intensities of C III X9710 than HD213409,

but all three were used in the mean. The results

clearly demonstrate the relatively large amount of

energy radiated in the emission lines even in this

wavelength region where the spectrum is not so rich

as in the ultraviolet. Furthermore the WC stars

are about twice as extravagant in general as the

WN's. The relatively low values of r for types WN4

and WN3 are due to the lack of any strong lines;

most of the expected lines of N IV and N V are in

the ultraviolet (near and far) . We can also make

use of Underhill's data for HD191765 (WN6) to estimate r for the regions A3100 to X4900 and X5300

to A6800. Again allowing 10% for weak lines we find

r = 0.712 and 0.266 respectively. If we allow

r = 0.250 for the wavelength regions not covered

here by definite measurements, we find that over the

range X3100 to XllOOO the emission lines contain

0.373 as much energy as the continuum. The value

for a typical WC star would be at least twice as large, judged by the appearance of their spectra.

When more precise measurements become available,

similar calculations can be performed for all WolfRayet stars. The implications of this result and

those of Table 4 are extremely important: there is

a tremendous amount of energy involved in the production of the emission lines, not only in the ac- tual transitions producing the lines but also in the

energy required to produce the high degree of ionization, and any theory must take this into account.

Again, aside from some applications of the Zanstra

mechanism (which assumes conditions obviously not

fulfilled in Wolf-Rayet stars) , no theoretical discussion has considered this aspect of the problem.

e . Variations in Line Intensities and Profiles

The greatest changes in the intensity and shape

of both emission lines and absorption components occur in the spectra of binary systems. However

small irregular changes have also been reported for

a number of single Wolf-Rayet stars. We will discuss the latter here and defer the binaries to Section III. H. Smith (1955) has observed variations

in the weaker lines of HD92740 and HD93131 (both

WN7) ; the usually conspicuous He II A4340 and A4200

practically vanish on occasion. Also the violetdisplaced absorption components of Si IV X4088 and

N V A4605 change erratically. Bappu (1951) has re128
ported similar variations for N V X4603, N III

A4640, He II A4541 and He I X4471 in HD191765,

along with apparent central reversals in He II X46 86

and He II A5411. He also reported intensity variations in HD192163, but his results have not been

confirmed by Underhill (1966) . Struve (1944) has

noted changes in the strengths of the absorption

components of N V A4603, A4620 and He II A4541 in

HD151932 (WN7) . All of these observations emphasize

the similarity between WN7 stars and Of stars, which

undergo changes in the lines of N III and N IV (Oke

1954)

The only other single Wolf-Rayet star to show

spectral variations is HD50896. Wilson (1948) noted

variable radial velocities and profiles of N IV

A4058. H. Smith (1955) also noted peculiar variations in absorption components and the presence of

central absorptions in most lines when N IV A4058 was bright and did not have a central absorption.

More recently Barbon, Bertola, Ciatti and Margini

(1965) have commented on the strong atmospheric

activity implied by the observed spectral changes.

However we must bear in mind that these variations

are very similar to those observed in binaries and

that HD50896 may yet prove to be a binary.

The evidence for intrinsic variability of

emission lines in single stars is very incomplete.

The meager evidence which is available suggests

some slight variability but nothing so drastic as the changes observed in other types of stars (e.g.,

T Tauri stars) . Variations observed incidently

with the photoelectric spectrum scanner are usually

of the order of 10% or less. No systematic study

of these variations has yet been undertaken.

III. THE EFFECT OF BINARY NATURE ON SPECTRA

A large fraction of Wolf-Rayet stars are binaries with an early-type companion. The binary

nature is detected by variable radial velocities,

by presence of stellar absorption features , and by

the general "drowning" of emission lines in the

continuum of a hot star. Examples of the last effect can be seen in Tables 3 and 4: compare for

example the intensity of He II A10124 in HD211853

(WN6 + BOI) to that in HD165688 (WN6) , 91 to 317!

Several general comments concerning the effects of

binary nature have already been made , so that I would now like to describe in some detail the

spectral changes observed in binaries.

The most studied Wolf-Rayet binary is the

129
eclipsing system V444 Cygni (HD193576, WN5 + 06,

P = 4.21 days), discovered by Gaposchkin (1941).

The light curve has been discussed by Kron and

Gordon (1943, 1950) and by Kopal and Shapley (1946).

The most detailed spectroscopic discussions are by

Wilson (1940, 1942) and by Miinch (1950). We will

concern ourselves chiefly with the spectroscopic

details discovered by Miinch and with additional

features observed photoelectrically by Kuhi (1968)

The major changes in the emission lines are

as follows: The He II A4686 has a symmetrical

profile at both elongations but is distinctly

asymmetrical to opposite sides during primary and

secondary eclipse. Figure 10 is a photoelectric

profile of the line at secondary eclipse (j3 = 0.0)

and at quadrature (j2f - 0.75). Sahade (1957) has

noted the appearance of a sharp emission feature

superimposed on the broader line during both

eclipses. He suggests that this new feature is

responsible for the asymmetry and that it arises

from gas streaming between the two stars. No such

feature is seen photoelectrically, but it may be

10000

5000-

0

5000

1—I—

I—I—I—I—I—I—I—I—I—I—

I—I—I—I—i—I—

Hen 4686

'•*•••••••••••,••••••• ••

<^-0.0

J I I \ I \ I I I I I I I L 4600 50 4700 50

Figure 10, Photoelectric line profile of He II

A4686 in V444 Cygni at secondary eclipse ((j) = 0.0)

and at elongation ((}) = 0.75). Note the distinct

asymmetry at (j) = 0.0; the asymmetry is in the opposite sense at (}) = 0.5, i.e., at primary eclipse.

130
wiped out by the low resolution (a 2 A exit slit)

The N III lines also undergo remarkable changes:

A4542 (blended with He II) appears stronger at

primary minimum than at secondary and is weakest

at elongations. On the other hand A 4640 is

stronger at secondary minimum than at primary.

Also very remarkable is the appearance of an absorption component (v = -340 km/sec) in N IV A 3483

only at primary minimum. It is not present outside

eclipse! Other lines of N IV do not show this

component (as might be expected since X3483 is the

only N IV line to show an absorption component in

other stars) ; instead X4058 is strongest at primary minimum and is more often distorted by an overlying fine absorption feature (which persists

for less than 6 hours) when the Wolf-Rayet star is

farthest away from the observer. The N V AA4603,

4619 lines weaken considerably at times of conjunction, and the accompanying violet absorption

edges also change markedly. They are strongest at

primary minimiam, weaker at elongation, and disappear

completely at secondary eclipse. The emission lines

of He I are very broad, and the absorption component

of X3889 is strongest around primary eclipse, but

its intensity is quite variable. Its variations

however bear no relation to the variations observed

in the other spectral features. Clearly it must be

formed at a fairly large distance from the star, as discussed earlier. Similar variations of He I X3889 have also been found in other binaries (e.g.,

Hiltner 1944, 1945). Munch tried to estimate the

intensity changes due to eclipse but did not suc- ceed (because of insufficient accuracy) , other

than to note that He II X4686 decreased by ^ 12%

at primary minimxam and N V A 4 60 3 weakened considerably at conjunctions. Hiltner (1950) also noted

the very peculiar behavior of He II A4686 in CQ Cep

(HD214419, WN7 + 07, P = 1.64 days): the light

curve in A4686 showed two maxima coinciding with con- junctions, and the intensity was greater by 4% at

primary minimum than at secondary. He also con- cluded that the intensity was intrinsically variable.

Kuhi (196 8) attempted to resolve the situation

by obtaining light curves through secondary eclipse

in all of the stronger emission lines. He planned

to use the 0-star as an occulting disk to cover up

successive regions of the Wolf-Rayet envelope and,

by comparing the shapes and depths of the eclipse

curves for different ions, to determine the stratification (if any) of the envelope. Thus if the

131
temperature decreased outwards, lines of N V would

undergo a much sharper and deeper eclipse than those

of N III and vice versa. The results, however, did

not reveal the stratification but only served to

point out the extreme complexity of the system. A

few such light curves are shown in Figures 11 to 15.

Several effects are immediately clear: (1) There

is no correlation between eclipse depth and ionization potential for a given ion. (2) Individual

Figure 11. Eclipse curves of He II lines at

secondary minimum of V444 Cygni . The phase 0 is

measured in hours from secondary minimum as determined from the continuum light curve. The intensity is that of the line only and has been normalized to 1.000 at (j) = +30.0. The mean observed

errors are ± 0 . 005 , ± 0.039 and ±0.014 for

AA4686, 5412, and 6562 respectively.

132
lines of the same ion do not have the same eclipse

curve. (3) Some lines (e.g., N III A.4540) undergo

extremely peculiar behavior not showing a true

eclipse curve at all. (4) Other lines (e.g., N III

X4100) have a very asymmetric light curve. In addition all lines undergo a decrease in intensity

at primary eclipse (Wolf-Rayet star in front) which

is often much larger than that at secondary eclipse

(e.g., N V X4609: 50% decrease at primary and 30%

at secondary) . The intrinsic variability of the

lines was also clearly demonstrated: N IV A7112

decreased by 30% over 21 days, N III X4540 increased

by 7%, and N V A4609 decreased by 11% in the same

1.0!

1.00

0.90-

0.80

1 1

1

1 1

1

1

1 1 1

1 1

1 T1 •

. . _. NHL

• X4640 '• o X4540 0 _

8

0 •

• •o //

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•

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• ••

\

1 1 1 1 1 ° 90 , 1 1 1 1 1 1. Figure 12, Eclipse curves of N III lines at

secondary minimum. Intensity units and phase as in

Figure 11. The mean observed errors for both lines

are ± 0.045. Note the peculiar behavior of A4540

133
time. This strongly points to the random nature of

the process responsible for the production of the

emission lines. The interpretation of these results

in terms of stratification effects proved to be impossible because of several additional complicating

factors implied by the data. Firstly a large fraction of the light in the emission lines comes from

the side of the Wolf-Rayet star facing the 0-star,

and there is a strong indication that much of this

is concentrated in a stream between the two stars.

Secondly, various lines (e.g. , He II A4686) were

found to increase in intensity as secondary mini- mum was approached, and some of this increase was interpreted as additional excitation produced by

Figure 13. Eclipse curve of N III A.4100. Units

as in Figure 11. The mean observed error is ^ ± 0.0 36. Note the asymmetry of the light curve.

134
the 0-star. Thirdly, the Wolf-Rayet envelope is

optically thick in many of the emission lines, so that additional non-geometrical effects are produced such as the appearance of absorption lines

and large decreases in intensity at primary eclipse

(e.g. , larger than can be explained just by occultation by the WN star) . Finally in such an envelope the electron scatterers act as secondary

sources of line emission and effectively enlarge

the area in which a particular emission line appears to be formed. Hence no insight can be gained

as to the original size of the N V A4609 emitting

region, and therefore the stratification question

0.80 —I—I—II—I—I—I—I

\—I—I—I—

Figure 14. Eclipse curve of N IV A7112. Intensity units and phase as in Figure 11. The mean observed error is + 0.014.

135
is still not settled. The final resolution seems

to require simultaneous photoelectric intensity

measurements and coude spectra obtained throughout

one epoch. The simultaneous measurements are nec- essary to disentangle geometric and physical effects; observations at one epoch are to avoid com- plications from the intrinsic variability of the

emission lines. Another system showing complex

behavior is Y2 Velorum, a southern WC star which

has not been studied very extensively. H. Smith

(1955) reports narrow central absorptions in most

of the emission lines from H I to 0 V. He conFigure 15. Eclipse curve of N V A4609. Intensity units and phase as in Figure 11. The mean observed error is ^ 0.035.

136
firmed Perrine's (1918) observations of the behavior of violet-displaced absorption components of

He I which appear and disappear on a time scale of

days. At times double absorption lines are visible.

The extremely curious thing is that the normal absorption component of He I X38 89 remained unaffected

by this behavior at a velocity of - 1320 km/sec com- pared to -480 to -960 km/sec for the other components. This behavior is quite rare and has been observed only in 1918 by Perrine and in 1953 by Smith.

No explanation has been offered. Bappu and Sinvhal

(1955) have investigated the behavior of He II

XA6560, 5411, He I X5857 and N IV X4058 in CQ Cephei

and find it similar to that already noted for He II

X4686. In addition Ganesh and Bappu (1968) have

described marked changes in line profiles of He II

A4686, X4200 and N IV A4058 in HD193928, HD186943,

and HD211853. Thus changes in emission line intensities and profiles seem to be a common characteristic of Wolf-Rayet binaries along with variations in the He I X38 89 absorption component.

We should also discuss the effects of a WolfRayet binary system on the 0-star. The effects on

the absorption lines (Munch 1950) have already been

discussed. One might expect the 0-type companions

to be somewhat peculiar, but Heals (19 34) stated

that they seemed to be no different than ordinary

0-and B-stars. If mass exchange does occur during

post-main sequence evolution, the line strengths

of carbon, nitrogren and oxygen may be somewhat

peculiar. This depends critically, of course, on what part of the Wolf-Rayet component is exchanged.

If the outer regions only take part, then the line

strengths of the 0-star should be normal. Preliminary results (Kuhi and Conti) indicate that they are indeed normal

IV. INTERMEDIATE OBJECTS

a, Wolf-Rayet Stars Showing Lines of

both Carbon and Nitrogen

A number of southern Wolf-Rayet stars described by H. Smith (1955) apparently show strong

lines of both carbon and nitrogen. HD62910 (WN6-C7)

shows the normal nitrogen spectrum together with

strong emission lines of C III A6735, A5696, A4326,

C IV X5805, C III-IV A4650 and O V A5590, A3760.

HD90657 (WN4-C + OB) shows C III A4650 and N III

137
X4640 of roughly equal strength. Finally HD104994

(WN3) is the strangest of all. It represents the

highest state of excitation observed in a WN star.

The lines of N V A4603, X4620 equal He II X4686 in

strength; other N V lines are extremely strong. In

addition 0 VI XA3811, 3834 are quite strong. This

is the only WN star for which they are seen; usually

they are conspicuous only in WC6 stars. Such un- usual stars (Smith has several other interesting

peculiar objects, but these three should illustrate

the point) do not fit into the simple picture of

two separate sequences but instead demand a very

specific explanation. Can they be explained by

special excitation conditions or by post-main se- quence evolution? No satisfactory answers have yet

appeared, but we should not forget these strange

stars when we comfortably extoll the virtues of one theory over another.

h* Stars Showing Wolf-Rayet Features and Nebular Lines

HD184 738 (WC8, Campbell's hydrogen-envelope

star) is the most famous of the very few stars in

this group. It has been extensively discussed by

Campbell (1918) , Stoy (1935) , Aller (1943) , Struve

and Swings (1940) , Swings and Jose (1950) , H. Smith

(1955) and Andrillat (1958). It shows a WC8 nucleus,

but the total intensities and line widths are much

less than for normal WC8 stars. The nebular lines

are extremely strong: [N II] A6580, 5755, [0 II]

A3727, [S III] X6312 are all present. Figure 16

shows the infrared spectrum of this object from

A6000 to XllOOO as obtained^with a photoelectric

spectrum scanner with a 10 A exit slit. The

strength of the Ha blend is fantastically large,

the central intensity being 80 times as strong as the adjacent continuum. The details of the Ha profile are shown in Figure 17. The forbidden lines

also dominate the infrared spectrum: [S II] AX6717,

6731, 10284, 10320, 10370, [S III] XX9069, 9532 and

[0 II] X7320. The He I X10830 line is also quite

strong. The weaker lines are due to C II, C III

and H I. There is some evidence for N II in the

star: it may be a contributor to the blend at A 8440.

Because of the nebular lines and surrounding nebula

this star has often been called a planetary nebula,

but it is not at all clear what it is. A related

object is HD167362, which also shows a strong nebular spectrum and a stellar spectrum similar to that

of HD184738. Again nitrogen is strong in the neb138
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6520 6540 6560 6580 6600

X Figure 17. Photoelectric line profile of the

X6560 complex in HD184738 obtained with a spectrum

scanner using a 2 A exit slit. The profile is a blend of nebular lines of H, [N II] and stellar

lines of C II.

ula but absent from the nucleus. Finally a most

peculiar object, NGC 6543, somewhat akin to the

above two stars , displays emission lines of carbon

and nitrogen in both the nebular and nuclear spectra. These objects may be the intermediate stage

between planetaries and Wolf-Rayet stars, but this

hypothesis leads to great difficulties with the masses of the central stars and those of normal

Wolf-Rayet stars.

140
o. Central Stars of Planetary Nebulae

Extremely little is known about the Wolf-Rayet

stars comprising the nuclei of planetary nebulae.

In fact work in progress by L. Smith and Aller suggests that the spectra of a majority of stars formerly called "Wolf-Rayet" do not resemble the spectra of normal Wolf-Rayet stars. Many of the nuclei

are Of stars which naturally have quite sharp

lines. A large number of them are peculiar, showing

broad emission bands near C III-IV X4650, He II

X4686 and 0 VI AX3811-3834 and differing in structure

from star to star. The other emission lines are very weak. The presence of 0 VI again implies con- ditions of very high excitation, much higher in fact

than those in normal Wolf-Rayet stars. Six stars

remain which are basically similar to classical

Wolf-Rayet stars. One of these (HD184738) has already been discussed under (b) . Smith and Aller

note two points about these six stars: (1) they have

consistently narrower lines, and (2) there is a greater tendency to find moderately strong lines of

both carbon and nitrogen than in ordinary Wolf-Rayet

stars. Thus there are still basic differences between these two groups which must be reconciled with

any theory.

The implied masses of the central stars are much smaller than those of noinnal Wolf-Rayet stars

('^^ 1 Mq instead of 10 Mq) . Hence we must conclude

that the Wolf-Rayet phenomenon is attributable to

the physical conditions in the atmosphere (but still

perhaps induced as a consequence of some stage of

stellar evolution) and is not the exclusive property

of a certain mass range of stars.

The objects discussed very briefly in these

sections emphasize the embarrassing fact already

pointed out that we really do not understand the

nature of Wolf-Rayet stars. I hope that this symposium will help alleviate this dismal state of

affairs.

141
REFERENCES

Aller, L. H. 1943, . J., 92, 135.

Andrillat, Y. 1952, C.R., 2J4, 2154.

Andrillat, Y. 1953, C,R,, 236., 51.

Andrillat, Y. 1957, Ann, d'ap. Suppl.^ No. 2.

Andrillat, Y. 1958 , C.R., 246., 1160.

Andrillat, Y. 1962 , C.R.^ 254., 64.

Barbon, R. , Bertola, F., Ciatt, F. , and Margoni , R. 1965, I.A.U. Inform. Bull. Var. Stars^ No. 109.

Bappu, M. K. V. 1951, A.J.^ 56, 120.

Bappu, M. K. V. 1968, Q.J.R.A.S., 9, 83.

Bappu, M. K. v., and Sinvhal, S. D. 1955, A.J.^ 60 , 152.

Beals, C. S. 1930 , P.D.A.O., 4, 271.

Bockasten, K. 1955, Ark. f. Fys., 9., 457.

Campbell, W. W. 1918, Pub. Liok Obs.^ H, 220.

Code, A. D., and Bless, R. C. 1964 , Ap . J., 139i, 787

Edlen, B."1956, Vistas in Astronomy ^ 2_, ed. A. Beer

(London: Pergamon Press), p. 1456.

Ganesh, K. S., and Bappu, M. K. V. 1968, Kodaikanal

Obs. Bull., No. 185.

Gaposchkin, S. 1941, Ap . J., 9^, 202.

Glad, S. 1953, Ark. f. Fys., 1_, 7. Hallin, R. 1966, Ark. f. Fys., 32, 201.

Hayes, D. S. 1967, Dissertation, University of

California, Los Angeles.

Hiltner, W. A. 1944, Ap . J., 99_, 273.

Hiltner, W. A. 1945, Ap . J., 101, 356.

Hiltner, W. A. 1950, Ap . J., 112, 477.

Hiltner, W. A., and Schild, R. E. 1966, Ap . J.,

143, 770.

Johnson, M. 1954, Observatory , 74, 124.

Kopal, Z., and Shapley, M. B. 1946, Ap. J., 104 , 160.

Kron, G. E., and Gordon, K. 1943, Ap . J., 9 7, 311.

Kron, G. E. , and Gordon, K. 1950, Ap . J., Ill, 454.

Kuhi, L. V. 1966a, Ap . J., 145, 715.

Kuhi, L. V. 1966b, Ap . J., 141, 753.

Kuhi, L. V. 1968, Ap . J., 152, 89.

Miller, F. D. 1954, Ap . J., 120, 265.

Morton, D. C. 1967, Ap . J. ^ TsU , 535.

Miinch, G. 1950 , Ap . J., 112 , 266.

Oke, J. B. 1954 , Ap . J., 120., 22.

Oke, J. B. 1964 , Ap . J., 1A0_, 689.

Perrine, C. D. 1918, Ap . J., 47, 52.

Plaskett, J. S. 1924, P.D.A.O., 2, 346.

Sahade, J. 1957, Mem. Soc. Roy. Soi. Liege, 20., 46.

Smith, H. J. 1955, Dissertation, Harvard University.

142
Smith, L. F. 1966, Dissertation, Australian National

University, Canberra.

Stoy, R. H. 1935, P.A.S.F,, 47, 162.

Struve, 0. 1944, A-p . J., 100 , 189.

Struve, O., and Swings, P. 1940, Pvoc. Nat. Acad.

Soi.^ 26^, 548,

Swings, P. 1942, A-p. J.^ 95, 112.

Swings, P., and Jose, P. D. 1950, Ap . J., Ill , 513.

Thomas, R. N. 1949, Ap . J., 109^, 500.

Underhill, A. B. 1957, Mem. Soo. Roy. Soi. Liege^

20, 14.

Underhill, A. B. 1959, F.B.A.O., 11, 209.

Underhill, A. B. 1962, Ap . J., 136, 2.

Underhill, A. B. 1966, The Eavly Type Stars (Dordrecht,

Holland: Reidel Publ. Co.).

Underhill, A. B. 1968a, B.A.N. , 19, 173.

Underhill, A. B. 1968b, Ann. Rev. Astr. and Ap .

6, 39.

Wilson, O. C. 1940, Ap . J., 91, 379.

Wilson, O. C. 1942, Ap . J., 95^, 402.

Wilson, O. C. 1948, T.A.S.P., 60, 383.

143
DISCUSSION

Chairman: Jorge Sahade

Sahade: In summarizing the spectral features

of WR stars, Kuhi has divided his material into a number of topics: (1) The possible overlap between

the WC and WN sequences. (2) The distribution of

energy in the continuum. (3) Line profiles. (4)

Correlations between the various characteristics

listed in items (1) to (3) . (5) The relative amounts

of energy radiated in lines and continuum. (6)

Variations in line intensity. (7) Influence of binary character. (8) Other objects exhibiting WR

and quasi-WR phenomena.

I suggest that starting with item (1) , we follow this sequence in our discussion. Let me re- mind you, however, of Kuhi ' s suggestion that with

the exception of C IV X5805, there are no strong

lines of C or O in WN stars, and that there are no even moderately strong lines of N in WC stars. The

only exceptions are several southern WR objects

that show both C and N lines. As Kuhi discussed

these exceptions in item (8) , I suggest we combine

items (1) and (8) in our discussion.

Underhill : First I would say that 0 V is definitely present in WN stars, and O IV is probable.

The spectrum of HD191765 provides a good example:

The He II A5411 line appears to lie at X5427, whereas other lines show no such shift. The He II line

must, therefore, be a blend, and when you look in

the multiplet tables, you find that indeed one of

the strongest multiplets of 0 V comes in just the

right position.

The same spectrum provides support for Kuhi's

argument that C IV is present in WN spectra. The

C IV lines at AA5801 and 5812 would account nicely

for the A5806 emission, and the alternative suggestion of N IV doesn't really compete. Whereas

Hallin actually predicted one multiplet of N IV

observed in this spectrum, he didn't even observe

the A5806 emission in his plasma source. It can,

therefore, hardly be strong enough to produce the

observed feature in the WR star.

Second, I am convinced that N III is definitely

present in the late WC subclasses. Kuhi has said

quite correctly that my evidence for N III in

HD192103 (WC7 or WC8) is not strong, but I say it

is not so weak that you can ignore it. The evi144
dence is stronger in Campbell's hydrogen-envelope

star, CWC9 on the present system) , where there is no alternative identification for A4634 but N III.

I would not suggest the presence of N in WC5 and

WC6 stars, but I think it is definitely present

in the later WC ' s

Alter: The forbidden lines of N II AA6548

and 6584 are strong in the planetary nebula sur- rounding Campbell's star, so is it really so sur- prising to find nitrogen in the stellar spectrum?

JJnderhill : It's no surprise to me, but people

insist on forcing nitrogen out of these stars when

it's got to be there. Kuhi suggested that because

the ionization potentials of C III and N III are nearly equal, it is difficult to understand why one should be present and the other not unless there are very real abundance differences. But in an emission

spectrum, the recombination spectrum is not necessarily governed by the ionization potential. Other

things such as recombination coefficients and the

question of what upper ions are present can also

effect the spectrum, so I don't think his remark

is terribly relevant.

Kuhi: I have already commented at length on these objects. I agree with Anne that WC stars

contain nitrogen. But perhaps we are belaboring

this point. The main point is that we have two

distinct sequences: one has strong lines of nitrogen, the other of carbon and oxygen. I do not

think it is particularly significant that we find

some nitrogen in stars of the carbon sequence.

Taken as a whole, the data on Campbell's star, on planetary nuclei and on the southern WR stars

lead me to believe that we are dealing with a "WR

phenomenon" , something which occurs because the

conditions in these objects are right for its oc- currence. I do not know what these conditions are,

but they do not appear to be uniquely related to a given type or mass of star. We find them occurring

in very massive and luminous objects similar to 0-

stars, and we find them in less massive objects

such as the nuclei of planetary nebulae.

Alter: I have a few remarks concerning the

nuclei of planetary nebulae, based on an observational program conducted by Lindsey Smith and myself. I want to emphasize that planetary nuclei

include a number of different spectral classes. Some

show only continuum; some, such as the nucleus of

NGC 6508, show absorption lines like 0-stars; some

are like Of stars, and some like WR stars. The distinction between the Of and WR stars is - or should

145
be - that those spectra showing narrow emission

lines are consistently called Of. The WR-type,

i.e., those with broad emission lines, can then be

divided into two distinct groups: those such as NGC

40 (Figures 18 and 19) and Campbell's star (Figure

20) whose spectra resemble the classical WC stars;

and those such as NGC 24 6 whose spectra are dominated by O VI AX3811 and 3834, C III-IV X4650, and

He II X4686.

I call this second group the 0 VI sequence be- cause its members show a continuous variation in the

strength and structure of the 0 VI lines (Figure

21) . At one extreme is the nucleus of NGC 246,

which shows a continuous spectrum with a few absorption lines and two sharp but faint emission lines

at A 3811 and A 3834. The O VI lines are somewhat

stronger in objects like the nucleus of NGC 2371-2 , which has faint broad emission lines topped by

sharp emission peaks. In IC 2003, the 0 VI lines

are still rather weak, while in IC 1747 they are just apparent, and in NGC 7026 they are prominent.

In NGC 6751 they are moderately intense, while in

NGC 6905, they are very strong and blended into a single band. Still more extreme than NGC 6905 are objects like NGC 5189 in which the O VI lines are enormously strong and completely dominate the spectrum. This star has been studied by Blanco, Kunkel

and Hiltner, who suggest that it may be the optical

counterpart of the x-ray source Centaurus XR-2. A

second object, with a similar spectrum but with no associated nebulosity, has been identified with

another x-ray source GX3+1.

It seems likely, therefore, that the O VI se- quence is one of increasing excitation and that the

most highly excited members may be strong x-ray

sources. I conclude that it is quite distinct from

any type of sequence found in the classical WR stars,

and I think it offers interesting possibilities for

speculation.

Sohild: Is there any correlation between the

properties of the stars in the 0 VI sequence and

such things as surface brightness and true diameter?

Alter: The statistics are so limited that it

is impossible to derive correlations. NGC 6905 and

NGC 2371-2, for example, are large nebulae, whereas

NGC 7026 and IC 1747 are binucleated. And while

the excitation level of the nebular spectrum is

often rather high, the 0 VI group does not include

the highest excitation nebulae. Even NGC 5189 does

not have very high excitation.

Johnson: Do you find any abundance differences

between the central stars and their associated neb146
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150
ulae? In particular, I am interested in the H

and He abundances : Do you find a higher He abundance in the nucleus than in the nebula?

Alter: I might be able to answer that if you

could tell me how to analyze the spectrum of a WR

star or planetary nucleus. For the nebulae, the

analysis is straightforward in principle, although

uncertainties in cross sections and computational

details make it more difficult in practice. For

the central stars, however, we know very little

about the formation of spectral lines, and in addition there is the problem of disentangling the

hydrogen spectrum of the star from that of the nebula. Current evidence suggests that there are no appreciable abundance differences between the planetary nebulae and their central stars. Perhaps the

best approach would be through the absorption lines,

since the absorption line objects seem to be very

much like ordinary 0-stars. However, Anne Underhill tells us that we can't even interpret these

spectra.

Wrubel: Is there any indication from their

position in the Galaxy that all these objects belong to the same population? Is the distribution

of these WR nuclei consistent with that of the

classical WR stars as outlined by Lindsey Smith?

Alter: Again the statistics are bad: we have

very few objects, and those that we have are very

faint. The low dispersion spectra were all taken

at the prime focus of the 120-inch telescope. Even

with electronic image converters, there is only a limited number of objects we can reach in the time

available. The answer to your question seems to

lie in more telescope time.

Smith: I have two comments on the differences

between planetary nuclei and classical WR stars.

First, the differences in line width mentioned by

Kuhi show up very clearly among the low-excitation

spectra (e.g. , BD+30°36 39 and NGC 40 as compared

with HD164270 and HD192103) , but in the high-excitation spectrum of the nucleus of NGC 6751, the

widths are comparable with those of classical WC6

spectra.

Second, the nuclei of planetary nebulae are believed to be of about one solar mass , whereas

the WR stars are about 10 M©. I have suggested

that the mechanism responsible for the WR spectrum

is an instability in the He- or C-burning core of

the star. The helium instability sets in only

above 7 or 8 M© , so if the same mechanism is to

account for the spectra of planetary nuclei, they

151
must have a C-burning core. Theoretical calculations suggest that a pure C-burning star is unstable

at any mass, but a hydrogen envelope would tend to

stabilize the system. So it seems reasonable to

suggest that the WR phenomenon in planetary nuclei

is due to C-burning instability.

Thomas: RR Lyrae stars are pulsationally un- stable and are of about one solar mass. How do

they fit into your generalization? If you assert

(1) that the evolutionary path of a star is fixed

by its initial mass and chemical composition, and

(2) that all stars within a certain range of these

parameters will go through a WR stage, how do you

interpret the difference between RR Lyrae stars

and planetary nuclei? Is it a difference in initial chemical composition? Or do they represent

different phases in the evolution of stars of the same initial mass and composition? Or again is it

possible .that they started with different masses

(and possibly different compositions) and that the

heavier one, following a different evolutionary

track, has lost mass, so that the present equality

in mass is a result rather than an initial condition? Both RR Lyrae stars and planetary nuclei are single stars, so it can't be a binary effect.

Smith: It is a different form of instability

in the two cases. For RR Lyrae stars, as for ce- pheid variables, it is K-instability due to the

variation of opacity with temperature and pressure.

For the WR stars, I am suggesting e-instability

,

due to the variation of energy production with temperature and pressure. This difference certainly

reflects a difference in the present structure of

the two classes of stars. How they attained this

structure is another matter. It is not, however,

clear to me that it is useful to compare these

quite different objects.

In this connection, I should like to add a postscript to my discussion yesterday on the distribution of WR stars and its possible interpretation in terms of initial chemical composition. I then made no distinction between binary and single

stars. Now you will recall from Schild's discussion that the association Sco OBI contains a WN7

star and a WC7 binary. Since the two stars are in

the same association, they must have essentially

the same age and initial chemical composition (although Schild did indicate there is some evidence

for small variations in these parameters) . The

fact that single WC7 stars do not appear in asso- ciations tells us they are fairly old. Yet here

152
we have at least one binary WC7 star that is definitely very young. I do not think we can escape

the conclusion that binary WC7 stars are younger

than single WC7 stars. If this is true for one subclass, it is reasonable to assume it is true for

all subclasses. Hence the binary stars and the

single stars are basically different and may have

different distributions in the Galaxy. While I still favor differences in initial chemical com- position as a working hypothesis, the matter does

need re-evaluation.

Thomas : Then let me continue the point I was making on instabilities. Suppose I assert that the

WR phenomenon reflects an atmospheric state (regardless of the mechanism by which that state is

produced) rather than a unique evolutionary stage

in a particular class of objects. And suppose further that this atmospheric state reflects a mechanical heating and that the variation in the WR phe- nomena reflects variations both in the character of

the heating and the character of the object (e.g.,

composition and gravity) . Then the kind, the degree, and the result of the instability must be

considered carefully. What distinguishes the WR

phase of the planetary nucleus, with its stable

broad emission lines, from the cepheid spectrum of

the RR Lyrae star, with its sporadic, narrow emission lines? Both these stars are of one solar

mass, and both have mechanical heating from radial

pulsation. Is it the "background" state of the

atmosphere, as determined by the size of the ra- diative flux? Is it the chemical composition of

the atmosphere? Is it the period of the mechanical

pulsation? Is it the total energy and momentum

carried in the mechanical pulsation? Superficially,

the WR instability appears to be violent, whereas

that of the RR Lyrae stars , the cepheids , the long

period variables and the Sun are mild. But the

terms "mild" and "violent" must be much more precisely defined.

J. Cox: 1 would like to comment on the physical distinction between the two kinds of instability. In the RR Lyrae stars and classical cepheids,

the instability is caused by ionization in the

stellar envelope. Since there is a limit to how

much an element can be ionized, there is a natural

limit to the amplitude of the pulsations driven by

this energy source. In the WR stars, on the other

hand, the pulsations are supposed to be driven by

nuclear energy sources in the core of the star.

Whether or not there exists a natural limitation to

153
the amplitude of such pulsations is not yet known.

It is a crucial point. One of our students, Ziebarth, is investigating this problem in connection

with the upper mass limit of the main sequence

stars, which owes its origin to the same kind of

instability mechanism.

Thomas: I was thinking not just of the amplitude of the radial pulsation, but also of the magnitude of the mechanical energy delivered to the

atmosphere. We know that in the Sun we get comparatively narrow, high-excitation emission lines in

the rocket UV. Some years ago Mrs. Gaposchkin

classified the quality of this spectrxam as WC6. We

have also heard Stecher comment on the great width

of emission lines in the rocket UV in C Puppis , and

he says this is also true of other supergiants.

Now suppose we were to increase by the same factor

the amounts of material and of mechanical heating

in the solar chromosphere-corona. Would we then get

broad emission lines similar to those in WR stars?

Or would we have to increase the supply of mechanical energy by more than we increase the atmospheric

mass? We have several parameters: energy available

(amplitude and period of pulsation) , energy delivered

(actually dissipated in the atmosphere) , excitation

level at some particular place in the atmosphere

(chromosphere-corona, observed only in rocket UV

or eclipse) , and excitation level in the great mass

of the atmosphere (disk spectriam observed in visual

region)

JJnderhill : The interesting thing about WR stars

is that they give no visible evidence of pulsation

in the sense that RR Lyrae stars and classical ce- pheids do. Kuhi has some evidence of erratic variations of small amplitude; but if there were pulsations , we would expect light variations and a period.

Thomas: Pulsations of the kind discussed by

Kippenhahn and Paczynski (as summarized by Lindsey)

have periods of minutes only. Would you detect

such short periods from the observations presently

available? Second, it is not the pulsation as such

but the mechanical energy which will affect our interpretation of the spectrum. If the pulsations

are of sufficiently high frequency that their effect is that of a statistically constant supply of

mechanical energy, it is not obvious that we would

observe a light variation.

Leung: On the siibject of the amplitudes and

variations in the radial pulsations of RR Lyrae

stars and classical cepheids, I would like to draw

154
your attention to an interesting correlation. For

classical cepheids, RR Lyrae stars, 6 Scuti and

dwarf cepheids, there seems to be a strong correlation between surface gravity "g" and the slope

"a" of the pulsation. We define "a" by Am =aA(RV),

where Am is the amplitude of the light variation,

and A(RV) is the amplitude of the radial velocity

variation. The larger the value of "g" , the smaller

the value of "a". If this relation were to hold for

WR stars, the very small amplitude of the light

variation would imply a small value of "a", which

would require a large value of "g".

Underhill : There is no evidence that WR stars

have particularly large values of "g". If you take

their average mass and Hanbury Brown's estimate of

their radius, you come up with log g 4.

Leung: Since we know nothing about the mechanism of instability in WR stars, I just wanted to

point out what would happen if you related them to

RR Lyrae stars and classical cepheids. Personally

I doubt that the mechanism is the same in WR stars

as in the others.

Sohild: I would like to make a point about

the separation of the WC stars into two groups.

Figure 24 from the Hiltner and Schild atlas [Fig- ures 24-28 appear at the end of the discussion,

pp. 175-179.] shows WC stars of higher excitation;

Figure 25, those of lower excitation. These objects were classified without any knowledge of the

identity of the stars. If you compare the strengths

and widths of the emission features around X4000,

you will see that they appear sharper and stronger

relative to the continuum in Figure 25. Whereas

Figure 24 consists of ordinary WR stars of WC type.

Figure 25 includes many peculiar objects such as the nucleus of NGC 40 and Campbell's hydrogen envelope star. These differences cannot be explained

simply as temperature or excitation effects. I just want to emphasize that wholly empirically there

are differences in appearance between the classical

WR objects and those associated with planetary nebulae .Sahade : We are all agreed that there are two

sequences of Wolf-Rayet stars. Kuhi and Miss Smith

argue strongly that they are differentiated by differences in chemical composition; Miss Underhill

suggested in 1958 that they could arise from differences in excitation. I should like to hear some discussion on this point.

Underhill : 1 think any such discussion should

be deferred until session C, when I have summarized

155
the various processes that must be taken into con- sideration.

Stephenson : As we are considering Kuhi ' s categories (1) and (8) together, let me suggest that

since old novae show broad lines. He II emission

and hot continua, they should be included among

"objects exhibiting quasi-WR phenomena".

Smith: If we believe that the characteristic

spectrum of the WR stars is due to excitation of

the outer envelope by a shock-wave, then we may

hypothesize that in novae and supernovae you get

similar characteristics because again you have

shock-waves running through the extended outer re- gions of the star.

Underhill : This epitomizes the fact that spectroscopic phenomena, as defined by spectroscopic

class, are not uniquely related to the parameters

basic to theories of stellar structure, i.e., mass,

composition and effective temperature. Our aim in

interpreting stellar spectra has been to derive

these parameters empirically from a knowledge of

spectral type. When we try to interpret B- and 0-

spectra, we find that somewhere around spectral

type B2 , we begin to run out of a unique relationship between spectral type and basic parameters.

The fact that we are now discussing quasi-WR phe- nomena as distinct from WR stars means that we have

completely run out. We have created a spectral

class that represents a physical situation in a plasma; the dominant characteristics of that situation are not uniquely related to mass, radius and

total radiation field. In other words, the WR spectrum alone cannot be used to establish the interior structure of the object producing it. We have

got to find some other feature in the stellar ra- diation field that can be used as a criterion of

structure

Alter: I wonder how many novae really fit into

this category of objects in which the continuum is

less highly excited than the line spectrum and in

which absorption lines appear on the violet edge of

emission bands. And furthermore, conventional no- vae are very shortlived phenomena, whereas the WR

objects endure for many years.

Underhill : We will have to define our criteria

more carefully. An important property of WR stars

is that they are steady. We could eliminate novae immediately by requiring that the star not change

rapidly with time.

Thomas: 1 am not sure you want to, Anne. I think Stephenson's emphasis on old novae was well

156
taken. When you start considering whether the two

sequences of WR stars reflect differences in composition or differences in excitation, you must re- member that Henry Smith (1955, Thesis) found one old nova that had changed from WN to WC or vice

versa. So I think we are demonstrating that while

we may have WR objects, we also have WR phenomena.

Lindsey has summarized one possible mechanism for

producing WR objects, namely that by which mass ex- change in close binaries results in an unstable Heor C-burning core. Possibly there are other ways to

produce WR objects. I think that old novae and

other objects exhibiting a quasi-WR spectrum may

give us a clue to the excitation problem as a guide

to the variety of ways in which a quasi-WR phenomenon - and thus perhaps the WR-phenomenon - may be

produced. They may also give us a clue to the origin of the two WR sequences

Regarding Anne's comments on the relation between stellar spectra and the parameters basic to

theories of stellar structure, I of course agree

enthusiastically. Twenty years of trying to develop a non-LTE diagnostic stellar spectroscopy have

been based on just this viewpoint and were inspired

to a large extent by emission line objects such as WR stars. But the first problem is to determine

what parameters are needed to describe the spectrum;

we cannot simply assume that we know what they are.

S. Gaposohkin : Let's not forget that many - possibly all - old novae are binaries. Perhaps the

analogy with WR objects or phenomena goes even

deeper.

Sahade: If there is no more discussion on this

point, I suggest we turn to Kuhi ' s second item,

namely the distribution of energy in the continuum.

Underhill: Kuhi's results present great problems. If we consider hydrogen-helium atmospheres,

we can predict the intensity distribution from models

computed by Mihalas, Strom, and myself. In the wavelength range we considered, you cannot change the

intensity distribution by going to higher temperatures; it is almost insensitive to the details of

the model. Yet the spectra which Kuhi has measured

and has corrected as well as possible for interstellar reddening show a significantly different

distribution. What can you do? You have no alternative but to postulate some unknown process.

Thomas: You agree wholeheartedly with our objections to the use of conventional models for the

interpretation of line spectra. Why don't you take

the next logical step and question the assumptions

157
underlying your continuum models? In addition to

those of LTE and radiative equilibrium, which you

have questioned, you have one' strong assumption

that you have not questioned - that of hydrostatic

equilibrium. You will remember the work of Kosirev

and Chandrasekhar during the thirties: They studied

the effect of arbitrary variations in density (such

as those that might result if hydrostatic equilibrium were dropped) on the distribution of spectral

energy. They happened to be looking for UV excesses,

but that is not the crucial point. I think we had

better stop talking about an atmosphere in hydrostatic and radiative equilibrium, characterized by

the two parameters "g" and T^ff.

Kuhi : Are there any models available for OB

supergiants?

Underhill : No. You can't get hydrostatic

equilibrium, and computers only work with hydrostatic equilibrium.

Kuhi: Okay, so that only emphasizes the point

I made yesterday about the similarity in the energy

distributions of OB supergiants and WR stars. And

it of course makes Dick's point all the more clear.

Steoher: Hunger and Klinglesmith have recently

analyzed the helium stars and found them to contain

40% C by mass. If the WR stars are overabundant in

C and He, then the C-absorption in the continuum

will change the atmospheric structure and the radiation field in the far-UV. In principle this suggests

a method for determining the C-abundance.

Underhill : Yes, to get a major change in the

intensity distribution, you do have to move away

from a predominantly H-atmosphere in hydrostatic

equilibrium. But I would not like to use the con- tinuous spectrum for an interpretation of the line

spectrum. I am sure we are observing two quite different plasmas. An interpretation of the plasma

responsible for the continuous spectrum is not likely

to be valid for the plasma in which the emission

lines are formed.

Sahade : We now move along to Kuhi ' s third

point: line profiles. He has shown us that most of

the profiles are Gaussian, but that some of them,

especially those with violet absorption edges, are flat- topped, and that many of those that are flattopped are just those lines that we would expect

to show the effects of diluted radiation.

Nariai : Do you mean by Gaussian profiles that

you expect turbulent motion of the order of 1000

km/sec? I am afraid such turbulent motion cannot

exist because it would have a decay time of the

158
order of 1 sec. You had better interpret the profiles in terms of a velocity gradient in the ex- pansion velocity. Flat-topped lines may be formed

in a region where the velocity is almost constant,

while Gaussian profiles originate in a region where

the expansion velocity changes rapidly.

Kuhi : I don't know that it is turbulent ve- locity, but there are two reasons for suggesting

that it might be. Firstly, the line profile is

roughly symmetric and roughly Gaussian. If I ignore

any radiation transfer problem, I find that random

velocities from 500 to 2000 km/sec are required to

give this shape. Secondly, where absorption com- ponents exist for such lines, their violet displacements are of the same size.

Thomas: In a cynical kind of way, I would re- mind you that just as most emission gradients appear to be exponential, so most line profiles appear to be Gaussian, at least until you start trying to fit them in self-consistent, numerical detail. You usually manage to find deviations once you know what you are looking for. The physical

problem in WR stars is to decide whether there is

indeed a random distribution of velocities at each

point in the atmosphere , or whether the atmosphere

is sufficiently transparent that the symmetric re- flection of a systematic velocity gradient gives the

impression of randomness. Clearly there are two

important physical questions, both of which we have

been trying to answer for the past 20 years. (1)

Can there exist random motions so violently supersonic as these velocities suggest? This was Nariai's

question, just as it was mine a long time ago when

I proposed an electron temperature of 10^ °K for a WR atmosphere. (2) What is the line opacity in the

atmosphere? So far we have been held up in answering these questions by our lack of a complete set

of solutions for the combined aerodynamical and

astrophysical treatment of energy dissipation and

radiation transfer in such an atmosphere. We have

been working on this piecemeal for twenty years.

A, Cox: Perhaps Kuhi would review the observational evidence for a "turbulent" shell with the

velocities just mentioned.

Kuhi: The evidence is indirect. Any simple

model of a WR atmosphere fails. If we assume pure

radial expansion, we find from eclipsing systems

that the size of the envelope is something like

5 times the stellar radius. But if the envelope

were that small, we would expect to observe

occultation effects, i.e., the emission lines

159
should appear asymmetric because the red-shifted

wing, formed in the receding part of the envelope,

would be hidden by the disk of the star. No such

asymmetries are observed. If we try forced rotational equatorial ejection, such as Limber has suggested, we again run into difficulties with profiles.

But if we assume that by some chance the envelope is really larger, we should then expect to find

phase differences between the time of eclipse as measured photometrically and the time as measured

from velocity curves. No such differences exist.

So we ask: How can we avoid these difficulties?

And one answer seems to be turbulence because then

any given line is the resultant of many components

across the disk of the star.

A, Cox: What about these lines that show

asymmetries at the eclipse phases?

Kuhi : That is a phenomenon associated directly with the fact that you have two stars and should

not be confused with the case of single stars.

Limber: As has been mentioned, turbulence with

such velocities is violently supersonic; the dissipation of energy would be enormous and would lead

to very great difficulties.

Underhill : It is misleading to use the word

turbulence in the sense that it is used by aero- dynamicists. What we really mean is a distribution

of velocities that appears from a great distance to

have the form exp(-v /b^). There is absolutely no doubt that the most probable velocity "b" is of the

order of 10^ km/sec. This corresponds to a temperature of some 10 ^°K, which is foolish.

Thomas : We ' re not saying a thing here about

turbulence that wasn't said twenty years ago. One

of the things we had hoped to get from this symposium was an answer to the question: Has there been

sufficient change in the phenomenological boundary

conditions or in our theoretical understanding of

aerodynamical and astrophysical problems to give us

a better understanding of the physical situation in

WR stars? So far in this symposium we have heard a few encouraging items: (1) The suggestion that

oscillations of a He- or C-core could provide a source of mechanical energy; (2) A good deal of

evidence that Tg is between 3xlO'*°K and 5x10 ""^K in

the region where the continuum is formed; (3) Confirmation that the radius of the region in which

the lines are formed is roughly a factor of five

larger than the radius as measured in the continuum. (4) Information on the rocket-UV spectrum,

which agrees well with results obtained in the visi160
ble; (5) Arguments in support of differences in

both chemical composition and excitation as an ex- planation for the existence of two sequences: WN

and WC. On the one hand, differences in taxonomic

properties are interpreted in terms of differences

in initial composition; on the other hand, there

is an overlap of N features in WC objects and of C

features in WN objects.

Underhill: Kuhi and I have been presenting you

with just such information on boundary conditions.

It is information that has existed for only a few

years and even then for only a few stars.

Let's turn to the flat-topped profiles. These

are not prominent in the ordinary WR stars. One of

the sources of confusion in the early days was the

happenstance that Beals chose to study C III X5696

in HD193793, which turns out to be about the best

flat- topped profile in existence. Naturally he said,

"expanding atmosphere". Other people, looking at

other profiles, reached other conclusions. There were two viewpoints in head-on collision, and both

seemed perfectly correct.

Beals: I think if you used the ends of the

flat- topped profiles to try to derive the velocity

distribution, you could get a Gaussian distribution

there too, although none of these interpretations is

very firm.

May I now ask why the violet absorption edges

are so strong in C III and He I, and whether they

are formed close in to the star or farther out? I presume you would say He I is farther out than He II

because it shows an absorption edge, which we attribute to dilution effects in the He I metastable

lines

Kuhi: You must be careful in your conclusion.

You might say that the He I metastable lines are formed farther out because of dilution, but you can't

take the next step and assume this implies Tg de- creases outwards. We don't know that.

Gebbie : Are the absorption edges confined to

any particular classes of WR stars, and if so is

there any correlation between these classes and the

classes which tend to produce small ring nebulae?

Smith: I believe the absorption edges are strongest and most consistently present in the spectra of stars in precisely those subclasses which

are associated with ring nebulae, i.e., the single

WN5, WN6 and WN8 stars. In the WC sequence, they occur most frequently in the WC9 spectra.

Thomas : Katharine and Lindsey are talking about

a correlation between absorption edges and nebulos161
ity; Kuhi has said that absorption edges are usually

associated with flat-topped profiles. Are we to

conclude that there is a correlation between nebulosity and flat-topped profiles?

Kuhi: What I said was: Absorption edges tend

to occur in lines which show flat-topped profiles,

if we see flat- topped profiles. I don't want to go

beyond that statement.

Underhill : What you see is a variety of com- binations. Flat- topped profiles with absorption

edges tend to show the characteristics of an ex- tended atmosphere. Actually they are Of lines.

There are also flat- topped profiles, such as C III

X5696, which do not generally have absorption edges.

Finally, there are lines with absorption edges which

do not have distinguishable flat- tops. I think this occurs when two lines at the same wavelength are blended.

Smith: The most conspicuous absorption line in

the spectra of WN5 , WN6 and WN8 stars is the violet

edge of He I A3889. Is that line consistently flattopped?

Underhill : It is flat-topped, but there is a rounded emission of He II right in its middle, so you can only see it when it's strong.

Payne - Gapo s ohkin : X5875 is flat-topped and

shows a nice absorption edge; and it doesn't have

a He II line on top of it.

Kuhi: There are several different processes

involved in producing these absorption lines.

There is one group, which includes XA3889 and 10830,

with a lower metastable level, and there are other

lines, about which Anne will speak in the next session, that arise from normal permitted levels.

Thomas: 1 am trying to be as simple minded as possible in order to see which simple minded arguments hold and which evaporate. In essence you and

Lindsey are saying: Deals' original suggestion of

a simple expanding atmosphere is pretty good for some stars, because on the one hand we see nebulosity which suggests an expanding atmosphere or ejected shell, and on the other hand we see absorption edges which suggest the same thing. The

question is whether these particular stars are ex- ceptional and, if not, why we do not see this kind

of double evidence of an expanding atmosphere in

all WR stars.

Sahade : I think we should remember that we probably have at least two kinds of envelopes in

WR binaries: the thick envelope which surrounds the

WR stars, and an expanding extended envelope. The

162
different profiles may be connected through this

model.

Now let us turn to Kuhi ' s fourth point, the

various correlations. First we consider the linebroadening mechanism. Kuhi has pointed out that

AA/X is roughly constant for lines of a given ion,

which suggests that Doppler broadening is the chief

mechanism. We have already discussed some aspects

of velocity fields. Kuhi also summarized the various suggestions about the effect of electron scattering. Is there any further discussion?

Underhill : From an empirical standpoint there

is little to discuss. The correlation "AA/A roughly

constant" is simply true. I intend to give an interpretive discussion in the next session. Briefly,

the point I shall make is that you must be very careful in choosing your lines and in assigning them to

specific parts of your model. That is why the interpretations just discussed - expanding atmospheres

and random velocity fields - can both be partially

true. Certain strong lines will have contributions

from many parts of the atmosphere ; certain other

lines will be formed only in shell-type conditions.

Beats: I am interested in the question of ac- celerated versus decelerated expansion. Do I understand that Kuhi thinks the atoms are decelerated

outward? If so, what is the mechanism? Is gravity

sufficient in these rather large envelopes, or do we need something like collisions between the atoms

and the surrounding envelope?

Kuhi : My interpretation depends, on whether you

assume the temperature increases or decreases outward . Steoher: Pikelner suggested some 20 years ago

that radiation pressure in the resonance lines might

levitate the atmosphere. One of his students has

applied it to WR stars and gotten reasonable electron densities in the outer atmosphere. Solomon

and Lucy have been applying it to mass loss from

supergiants. I believe the ionization will increase

outward due to the decrease in electron density.

Thomas: The idea that radiation pressure in resonance lines can act to drag out the atmosphere

goes back to Milne and his suggestion about Ca"*" in

the solar atmosphere. I am not sure that it has ever been successfully embodied in a wholly selfconsistent theory. And it is not obvious that the

ionization will increase with an outward decrease

in electron density. The coronal type of collisional ionization equilibrium is independent of

density. I agree with Anne that many of these ques163
tions can be answered only in the context of a completely self-consistent model. All these jigsaw

pieces must be put together into a complete mosaic.

Sahade : Kuhi ' s second correlation is that between line width and spectral type in WC stars : the

broader the lines, the earlier the spectral type.

Kuhi: This is not really discussible; the line

widths are used to define the WC subclasses.

Sahade: Then we go on to the third correlation:

a decrease in line width with increasing ionization

level.

Thomas: This goes back to Beals ' first work

and is one of the oldest correlations in the literature. I took it literally in my model of the WR

atmosphere as an extended chromosphere-corona supported by mechanical heating. In such a model you

would expect T^, and hence the ionization, to in- crease outward, at least initially. Then the observed correlation would suggest an outward decrease

in the expansion velocity, or, depending upon how

you interpret the line broadening, in velocity gradient. But whatever the interpretation, I regard it

as one of the most significant pieces of empirical

evidence, or boundary conditions, that we have. I was glad to hear Kuhi reaffirm it.

Underhill : I don't believe it has been established observationally . It is not based on enough

quantitative information. I think it is just a happenstance

Thomas: Happenstance means it exists, but you

don't understand why. And here the important question is whether it exists, regardless of why.

Kuhi: Well, I think it is true in some cases. There are certainly spectra in which the lines of

N III, N IV and N V are really quite different.

But as Anne has pointed out, there are exceptions.

I am afraid this will always be our problem with

WR stars. There are exceptions to almost anything

you can find.

Underhill : I think it was Mrs. Gaposchkin who

first remarked that classifying WR stars is a pointless task: you end up with one class, one star.

According to Lindsey Smith, we have 124 WR objects

in our galaxy and 24 different classes of objects.

That gives you about 5 objects per class. I'm dead

against classification atlases: every time I take a high-dispersion plate, I find that some line which

has been described as, say, N IV or C IV, is a general muck of 10 or 15 lines. Empirical classifications may be perfectly consistent, but I don't

164
think they are accurate enough to establish this

sort of correlation.

Sahade: So we go on to the fourth correlation,

that between sharp-line objects in the WN sequence

(Hiltner and Schild class A) and binary stars.

Thomas: Can you be more precise? Is the difference between sharp and broad lines a matter of

10%, a factor of 2, or a factor of 5?

Kuhi : Possibly a factor of 2; nowhere near a factor of 5.

Vnderhill : This is one of the points that

requires further investigation and quantitative

measurement. Just how sharp are these lines? What

is the real meaning of conclusions based on lowdispersion spectra? But it is a very difficult

observational problem to get 20 &/mm spectra of

these objects.

Payne-Gaposohkin : But you do not deny the scientific use of the division; it is so wonderfully

physical. All sharp lines happen to be binaries.

When people saw this for the first time, they saw something important.

Roman: I'm somewhat confused on looking at the

Hiltner and Schild Atlas. [The Hiltner and Schild

Atlas, Figures 24-28, appears at the end of the

discussion, pp. 175-179.] Only about half the stars

in the sharp-line WN sequence (WN-A) (Figure 26)

appear to be binaries. In general, the single

stars in this sequence seem to have broader lines

than the binaries, but nothing like as broad as those in the broad-line sequence (Figure 27) , one

of which is in fact a binary. Then of the two WN8's

the single star appears to have slightly sharper

lines than the binary. So I'm a bit confused as to

the facts.

Underhill : But that's just the problem. When

you compare lines of a binary with those of a single

star, the width can be misleading. You should in

fact compare the half-widths not of the actual but

of the normalized profiles.

Schild: As I said before, the spectrograms

were arranged in sequences on the basis of their

appearance. It was then observed that a large

number of objects in one sequence were binaries.

That is all there is to say.

Sahade: Let us now go on to Kuhi ' s fifth item:

the relative amounts of energy radiated in the lines

and in the continuum.

Gehbie: In view of the large amount of energy

in the emission lines, I am curious about the phys165
ical significance of a visual magnitude that ex- plicitly excludes these lines. Could this account,

in part, for the anti-correlation between excitation

and luminosity obtained by Lindsey?

Underhill : If we assume this energy is a form

of the UV radiation, our bolometric correction will

account for it just as for a nomal star. The

strength of the emission lines can perhaps be re- garded as a conversion of the bolometric correction.

Gebbie: But these are not necessarily recom- bination lines.

Undevhill : Some are; some aren't. Those that

are, represent a good fraction of the conversion.

Thomas: Not obvious at all. You are assiaming

a mechanism. Suppose we buy the following picture:

a continuum corresponding to about S^IO'^^K formed

in a photosphere, and emission lines formed in a random-velocity shell heated by mechanical energy.

We then have two kinds of energy supply: the con- tinuum will refer to radiative processes, and the

lines to collisional. So the continuum absolute

magnitude and line absolute magnitude will refer to

two different processes.

Underhill : I agree. To get a meaningful total

flux, or bolometric magnitude, you should include

both the mechanical and radiative energy. The term

"effective temperature", as it is usually used, is

misleading. The mechanical flux is usually neglected, as its contribution is small compared with the

radiative. How much of the energy in the lines comes from the far-UV radiation field, and how much

from the mechanical flux is, I agree, an open question.

Gehhie: Then is it meaningful to exclude from

the absolute magnitude a source which may contain

half the energy?

Smith: My change in luminosity with subclass

amounted to some two magnitudes between WN3 and

WN8. Diane Pyper gives a mean correction for the

influence of the emission lines on the magnitudes

of WN stars that amounts to 0.2 mag. Kuhi ' s data

would give corrections amounting to between 0.1

and 0.3 mag. Thus the two corrections are in good

agreement and would be an order of magnitude too

small to account for the anti-correlation between

excitation and luminosity.

Kuhi: Let me emphasize that the values I gave

in my summary from my own measurements covered only

the region AX8000-11000 . To derive the ratios 37%

for WN stars and 70% for WC stars, I combined my

data with that of Anne Underhill.

166
Underhill : These corrections are only for My,

whereas it is the bolometric magnitude that is fundamental to theories of stellar structure. Implicitly we assume the bolometric correction for WR stars

is about the same as for 0-stars. But even if it

is, we still don't know the correction for 0-stars

to within half a magnitude. And if the emission

lines are excited by mechanical Energy, we may be

off by another half a magnitude or more. On the

other hand, we must realize that while lines may

produce 70% of the energy in the visible region,

that is still only a small percentage of the total

energy emitted by the WR star down to or below

912 A. We are looking at these stars in the faintest part of their continua.

Thomas: Could I ask if the following is a fair summary of what you are saying: We observe My

for the continuum. First, we want to correct it to

include the energy emitted in the lines , so that we

can estimate the total radiation in the visible region. Then, from this small visible tail, we want

to infer the size of the dog - the bolometric magnitude. Your bolometric correction is 2-3 mag,

Lindsey?

Smith: Less than 4.6 mag for Kippenhahn ' s Heburning stars; about 2.5 mag for the same stars in

the C-burning phase.

Thomas: So you are really saying that the bolometric corrections are so large and so insensitive

to the details of the model that no matter what

model you use - including just a blackbody - the

introduction of a supply of mechanical energy can have no significant effect. Right? Then this is an assertion which is basic to our entire discussion and which must be checked. I remind you the

WR star is likely to be the most extreme freak of

all freaks, and it is not impossible that it may

deviate from this simple assumption.

Undevhill : You are right in your summary of

what we do; but I don't think the situation is so weird when you begin to look at it; it behaves like

physics should.

Johnson: We've heard that a large fraction of

the energy is radiated in the emission lines and

that in binaries, a large part of the line-producing

region lies between the stars. Is it possible that

this same region produces a reasonable part of the

continuum by free-free emission?

Kuhi : My remark that a large part of the lineproducing region lies between the stars referred only

to V444 Cygni, whereas the figures on the energy ra167
diated in the emission lines referred only to single

stars

Sahade: Now when we turn to Kuhi ' s sixth point,

the variation in line intensities, we must distinguish between lines in single stars and those in

binaries. The latter should be included in item

(7) . Aller: I want to mention HD45166, because it

is usually quoted in connection with intrinsic variations in line intensities in WR stars. Apart from

this star, most discussions concentrate on Of stars

where strong variations in line intensities are well established. Using Harvard plates many years

ago, Carol Anger Piene found tremendous variations

in the nitrogen lines near AX4634-4640. The spectrum of something like a B-star, with well-marked

hydrogen lines, is also apparent. The star is

evidently a binary, but there is no simple periodicity. There is still some question whether

it is really a WR or an Of star. (Personally, I don't think we have yet settled the question of

the dividing line between the two.) The star re- sembles an Of star in the sharpness of its lines,

but the character of the spectrum is that of a WR

star. I do not believe these variations in the

nitrogen lines have anything to do with the binary

character

Sahade: If there is nothing more on non-binary

effects, we will turn to the binaries, item (7).

Smith: Kuhi found that in V444 Cygni the sec- ondary minimum (WR star eclipsed) is deeper in the

red than in the photographic. This may be due to

the infrared excess of the WR star, but how would

you interpret Hiltner's result (1950, Ap . J.^ 112 , 477) that in CQ Cephei we observe the opposite effect, i.e., the secondary minimum is deeper in the

UV than in the visual?

Underhill : Those were broad-band observations,

1000 A or more wide.

Kuhi: Then we can't say anything definite;

broad-band observations include the effects of emission lines. If we are to compare CQ Cephei and

V444 Cygni, we need narrow-band work.

Eq ellming : I'd like to comment on a WR eclipsing system that I think has been much neglected.

This is HD168206, CV Ser, found by Gaposchkin many

years ago. The system has a visual magnitude of

9.14 and a spectral type, according to Smith's

classification, of WC8 + BO. It is the only e- clipsing system known in the WC sequence. Figure

22 shows three light curves. The upper curve was

168
_ J. „ L 00

_ 1 - 0.2

_ I . 04 0.6 08

Phase

1.0 1.2 1.4 Figure 22. Photometric observations of CV Ser.

The upper curve was obtained by Gaposchkin (1949) , and the' lower curves are new observations. Note

the depth and width asymmetry of the primary eclipse,

The period is 29.640 days.

presented by Gaposchkin when he first showed it was

an eclipsing system: from this data it has a 0.14

mag primary eclipse and a 0.08 mag secondary eclipse.

About 5 years ago, in the course of a photometric

survey of binary characteristics in WR stars, I re-observed this system for about 6 weeks. The low- er two curves show the results of my broad-band UBV

photometry. The middle curve shows the variation

in the yellow; the lower curve, the results in the

blue. There are several points of interest. First,

the primary eclipse (WR star in front) is much

deeper than that observed by Gaposchkin: the blue

shows a depth of 0.55 mag, and the yellow, 0.49 mag.

At first we attributed the discrepancy to the fact

that Gaposchkin had used a slightly erroneous period

furnished by Hiltner. However, there is a faint

possibility that the envelope has actually changed,

that the primary eclipse is deeper because the WR

envelope is bigger. The second point of interest

is the difference in the depth of the eclipse -in the two colors: the blue eclipse is deeper than

the yellow by about 0.06 mag, which is well above

169
the errors of the photometry. Studies of the variation in emission lines during eclipses would be of

great interest in connection with this system. The

system has a period of 29.640 days and takes about

a week to go through primary eclipse, so detailed

spectroscopic studies could be done in leisure.

Finally, if you wonder why I have no data on the

secondary eclipse, it is presumably because it oc- curred during a period of bad weather.

Kuhi : First, although bright, the system is

still a magnitude fainter than V444 Cygni, so the

observing time required for a given accuracy is

greater. Second, the star has been observed photoelectrically at Lick in the hope of detecting the

secondary eclipse which is, of course, the one that

should be observed for changes in the intensity of

emission lines. No secondary eclipse has been detected to within the accuracy of the UBV photometry,

i.e., to, something like 0.01 mag. It is still possible you might find an eclipse in some of the emission lines, and it would be especially worthwhile

to pursue this. However, as we do not know the ex- act date of the secondary to within two days , I would hesitate to ask for time on a large telescope.

Perhaps it would be worthwhile to try a smaller

telescope, using narrow band filters (e.g., on X4686

of He II) and integrating over longer periods of

time. Then if something happens, we can go to the

120-inch.

5. Gaposchkin : 1 am delighted to see you are working on this system. As you mentioned, it is one of the few examples of a WC eclipsing binary.

So if we can tie down this secondary minimum, we

can determine the sizes and m.asses of the two com- ponents. I personally find you spectroscopists

the wildest and most entertaining group of all the

astronomers. I think this is because spectra are enigmatic phenomena. Struve once told me that he can take one spectral plate, work on it for an en- tire year, and get something out of it. This con- trasts with my own need of 1000 plates to get the

right answer for an eclipsing system. One plate

gives the spectroscopist a year's work; 1000 plates

give me one set of numbers.

Regarding 'Kuhi ' s statement that his observations of binaries are a mess, I think it is a most

revealing mess. He does great injustice to his

work. I think he can be the first to give a real

structure, in a graphic way, to a real WR star.

Figure 23 shows an exact dimensional picture of

V444 Cygni; the inner structure of the star is

170
V444 CYGNI

0 Component

Moss: 25©

Radius: 10©

Abs. mog.: \- 32000«K

WR Component

Moss: IOq

Radius: 6©

Abs. mag.: -z"!? \ : 14000 "K

Figure 23. The best and the only unambiguously

determined WR + O system as illustrated by Sergie

Gaposchkin.

schematic; the eclipse light curve is below. Kuhi

,

you should be able to use your spectroscopic data

to come up with an equally precise picture of the

structure

Kuhi : The real problem, however, is that I do

not have geometrical effects alone to consider.

There are all those other things I mentioned in my

talk that confuse the situation. I wish it were

possible to interpret eclipse profiles in terms of

geometry alone.

Sahade: In considering the WR binaries, we

are discussing two main problems. One is connected

with the He I lines, which show dilution effects;

the other is connected with the emission-line shape.

There is the strong evidence that the asymmetries

of structure observed in emission lines (e.g.. He II

X46 86) come from material located between the two

stars. This structure has been observed several

times in Y2 Vel, the first time by Perrine.

Hanbury Brown: The measurements of Bappu and

Ganesh, which show an orbital velocity ranging from

+200km/sec to -lOOkm/sec and a period of 78.5 days, were taken in the C III-IV emission complex at

171
A 4652. We measured the apparent angular size of the

system in this same line and found it to be about 5 times the size of the star. We therefore appear to

have a region 5 times the stellar diameter orbiting

in a binary system. Is this regarded as a reasonable model?

Underhill : The whole star is orbiting at this

speed. The extended atmosphere moves with the star.

That is perfectly normal for an early type binary.

Thomas: When Anne says this is perfectly normal, she means it is normal by assumption. In all

such discussions, including those on mass exchange,

the detailed aerodynamics of the problem are generally neglected. Clearly, there is an interesting

set of problems to be solved. Possibly we are being

unduly optimistic in supposing we can interpret the

spectral phenomena without simultaneously solving

the aerodynamic problem.

Sahade : Let me remind you of Kuhi ' s remark

that He I A3888 is always observed as a strong line

and is always greatly displaced - by the same amount

- to the ultraviolet. In stars like Y2 Velorum, the

line has several components less displaced to the

violet, the position and number of which undergo

very pronounced changes. Sky and Telescope of 1956-

5 7 contains a reproduction of consecutive plates of

Y2 Velorum, and you can see how fast these changes

occur. A similar thing happens in the well-known

peculiar system of 3 Lyrae , which is also surrounded

by a large expanding envelope, as indicated by the

presence of a violet-displaced line of He I X3888,

again with several less-displaced components.

Regarding the He II X4686 emission, Velorum

also displays a structure which may be connected,

as in V444 Cygni, with matter streaming toward the

companion star. The structure of X4686 in V444 Cygni

is not simple. Sometimes the relatively narrow,

superimposed emission looks double, and sometimes

there is even a sharp absorption cutting in. In

Y2 Velorum, the feature is even more complicated.

Kuhi: I must say that in V444 Cygni nothing

really correlated with anything; once you tried to

correct for secondary effects, there were no clearcut correlations. There was no clear cut correlation between eclipse curves (shape and depth) and

ionization potential. Individual lines from different ions behaved differently, as did different

lines from the same ion. I think the situation

with CQ Cephei is worse. Indeed, I think CQ Cephei

will probably confuse the interpretation of WR stars more than it will solve it because it is a contact

172
binary and we are dealing with the interaction between two stars.

Underhill : The purpose of this symposium is

to try to establish a body of observational facts on which to base our theory and interpretation. So

although these complicated systems can be used for masses and radii, I think we should put them aside

and concentrate on a select group of quiet, wellbehaved objects.

Sahade: Yes, but what if all WR stars turn

out to be binaries? At this point we are discussing

binaries; if they happen to be complicated, I can- not help that. But to judge from Kuhi ' s comment,

I think we have exhausted this subject.

Schmidt-Kaler : . As a postscript to this and to

Lindsey's summary in section A, I would like to

suggest a possible connection between nebulosities

and a new type of stellar aggregate which we have

discovered and called "stellar rings" [Veroff.

Bochum^ No. 1, in press]. A stellar ring is a cluster of stars which appears as a regular elliptical

ring with a very sharp outer boundary. The thickness of the ring is about 1/30 the minor diameter.

The number of stars involved averages 70 and may go

up to 200, so the density is considerably higher

than that of the general stellar field. For six

objects with photometric distance moduli, we obtain

a unique minor diameter of 7.1 psc and ages of between 0.5 and 5x10^ years. We believe these rings

must be the result of an expansion process.

In searching the Palomar Sky Survey prints for

the precursors of stellar rings, we found a number

of emission objects that display the same characteristics as the rings, except that they are gaseous instead of stellar. Their properties are summarized in Table 6. One might interpret these

data as suggesting that these nebulae represent a continuous transition from the simple ionization

front around a WR star, to the appearance of a shock

front, to the formation of. stars in a stellar ring,

and finally to the decline of the ionization, leaving only the longer-lived stellar ring. Thus I would suggest that the WR stars act as a kind of

"blasting cap" for star formation.

Johnson: The Russian astronomer Dolidze has

found that emission-line stars are concentrated

around supernova remnants - IC 4 43 was mentioned.

He applied similar reasoning to the production of

stars in compressed gases outside the expanding

shell, only here it was a supernova that did the

triggering.

173
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8901^22 N179

PART C

SPECTROSCOPIC DIAGNOSTICS,

INTERPRETATION, AND ATMOSPHERIC MODELS

CHAIRMAN: Lawrence H. Alter

Department of Astronomy j University

of California at Los Angeles

INTRODUCTORY SPEAKER: Anne B. Underhill

Sonnenborgh Observatory ^ Utrecht

I. Introduction 18d

II. The Types of Spectroscopic Information 184

a. The Continuous Spectrum 184

b. The Line Spectrum 185

III. Spectroscopic Diagnostics

and Interpretation 19 2 a. Evidence Concerning the Temperatures

of Wolf-Rayet Stars 19

b. Evidence Concerning the Electron Density

in Wolf-Rayet Atmospheres 19 5 c. The Types of Physical Processes

to be Considered 19 6 IV. A Model of a Typical Wolf-Rayet

Atmosphere 19 8 V. Clues Regarding the Evolutionary

Stage of Wolf-Rayet Stars 19 9 References 201

DISCUSSION 202

181

I. INTRODUCTION

The Wolf-Rayet stars have been selected as a group according to certain conspicuous characteristics of their spectra, which consist of a faint

continuous spectrum crossed by rather strong, broad

emission lines from He I, He II, C III, C IV, N III,

N IV, N V, O III, O IV and O V. Other emission lines

are present, but the most characteristic lines upon

which the spectral classification system is based

belong to these spectra. A few particular absorption lines are seen. The absorption lines are always displaced towards shorter wavelengths with

respect to the emission lines, which usually have

wavelengths consistent with the expected radialvelocity component of the peculiar motion of the

star. In this discussion variations in the spectrum due to orbital motion in a binary system and

to blending with the spectrum of a companion will

be considered to be of secondary importance. They

offer merely an extra source of confusion which

should be avoided so far as possible by a careful

selection of objects for study.

An attempt will be made to show what the observed spectroscopic details imply concerning the

physical conditions in the atmospheres of Wolf-Rayet

stars. Of themselves these conclusions say nothing

about the evolutionary stage of Wolf-Rayet stars.

However, when the information about atmospheric con- ditions is coupled with other astronomical information about Wolf-Rayet stars and about other stars,

it should be possible to draw a self-consistent picture of the position of Wolf-Rayet stars in the general pattern of star formation and evolution. The

present discussion will conclude with a few remarks

on this subject.

Spectral classes give an empirically arranged

ordering of stars according to changes of a few

characteristic features in the spectra of the stars.

One may not assume without further investigation

that the selected criteria lead to classes arranged

monotonically according to a theoretically significant parameter such- as the effective temperature

of the star. One of the goals of the present discussion is to see if it is possible to find criteria which do arrange Wolf-Rayet stars in a sequence

according to effective temperature. Effective temperature is selected as basic because it is a sig183
nificant parameter in the theories of stellar structure and evolution.

To go from observed spectral detail to effective temperature implies possession of a satisfactory theoretical understanding of the formation

of stellar spectra. We do not have this understanding at present. Much of what will be presented here is an attempt to formulate the chief

problems in understanding Wolf-Rayet spectra. More

questions arise than can be answered. However,

progress will have been made if we can establish

some of the significant questions.

In Section II a brief summary is given of the

relevant spectroscopic information. Full references

for this material can be found in other papers presented at this conference and in Underhill (1968)

The chief aspects of the spectroscopic diagnostic

process are discussed in Section III, while in Section IV -a qualitative model of a Wolf-Rayet atmosphere is presented. Clues concerning the evolutionary stage of the Wolf-Rayet stars are discussed

in Section V.

II. THE TYPES OF SPECTROSCOPIC

INFORMATION

a. The Continuous Spectrum

In comparison to 0-stars the WN stars have a definite ultraviolet excess which begins near 4000

A; the WC stars have only a slight ultraviolet excess, if any. Through the green-blue-violet spectral region the relative energy distribution in the

continuous spectrum of Wolf-Rayet stars is very

similar to that of 0-stars. Both WN and WC stars seem to have a small infrared excess. It is difficult to say how large this excess is because the

correction for interstellar reddening is a critical

factor in determining the true infrared excess. From a study of eclipse depths in different

wavelengths at secondary minimum Kuhi (1968) has

shown that the Wolf-Rayet component of V4 44 Cygni

definitely brightens in the infrared with respect

to its 0-type companion. This binary system con- tains gas streams. One cannot say surely whether

the brightening in the infrared is associated with

the Wolf-Rayet star itself or with gas in the binary

system lying in the neighborhood of the Wolf-Rayet

star. Such gas would be expected to radiate a free-free continuum in the infrared.

184
b . The Line Speotrwn

It is a significant fact that the emission and

absorption lines of Wolf-Rayet spectra cannot be

regarded as being formed in one body of gas at one representative temperature and pressure. Rather

one must acknowledge that the observed spectrum is

a composite of features which are formed under quite

different circumstances. Progress in understanding

the meaning of Wolf-Rayet spectra is dependent on the correct grouping of lines together for interpretive purposes. Successful analysis is hindered by the fact that we have only a limited spectral range accessible for study. The available information can be conveniently grouped under four

headings

1. Shapes of the emission lines

(i) There are a few flat-topped emission

lines which in some cases are accompanied by a shortward-displaced absorption component. These

line shapes can be explained quite satisfactorily

by the classical theory of a spherical expanding

atmosphere. In WC stars the line C III X5696 gives

an outstanding example of a flat-topped profile not

accompanied by an absorption component. In WN stars,

and in some WC stars, the lines He I X5876 and

X3888 have a flat-topped emission profile accom- panied by a strong shortward-displaced absorption

component. In WN6 stars the line N IV A 4057 appears to be broad and flat-topped in emission, but

there is no clear absorption component. On the

other hand the rounded emission line from the multiplet N IV A3478-82 is usually accompanied by a strong shortward-displaced absorption component.

In WC stars the rounded emission line from the C III

A4650 multiplet behaves in a similar manner to N IV

X3478-82. One cannot isolate a flat- topped emission

line for either of these multiplets.

The lines mentioned above are the chief lines

in Wolf-Rayet spectra that have profiles that can

be explained, at least in part, by the simple hypothesis of a spherical expanding envelope. It is

worthy of note that most of these lines appear in

emission in Of stars as a result of particular

processes or are known to be strengthened in absorption by dilution effects.

The spectrum profil§ of HD193793 (WC6 + O) in

the region 5630 to 5960 A* is shown in Figure 1.

The flat-topped profile of C III X5696 is clearly

185
5630 40 SO 60 TO 80 90 5700 X3 20 XI 40 50

1 . 1 . 1 . 1 . 1 . 1 . 1 , N

1 ' 1 ' 1 ' 1 ' 1 ' 1 ' 1 1.1,1,1,1,

/ HDI93793 /

t- /

/

1

C]

1

I

1 1 1 1 1 1 1

40 StJ-v ^.JBD 70 80 90 58C > 1.1,1,1,1 0 10 20 X) 40 50

Interstellar

_

NaD2

Interstellar

50 60 70 80 90 5900 lO 20 30 40 50 Figure 1. The spectrum profile of HD193793,

WC6 + O, between 5630 and 5960 A.

evident and also the flat-topped profile and displaced absorption of He I A5876. The C IV lines are discussed below. The spectrum of this binary

star has been chosen to illustrate the features

formed in the expanding atmosphere because the WolfRayet spectrum is seen blended with the essentially

continuous 0-type spectrum. Consequently the intensity of the emission lines is nowhere so great

with respect to the continuous spectrum that the

photographic photometry is seriously in error. The

flat tops are real.

The flat-topped lines in the spectrum of

HD191765, WN6, may be seen in Figures 2 and 3. The

flat character of the He I A5876 emission is striking; that of the weak emission due to He I A3888 is

less conspicuous, although the shortward-displaced

absorption core is strong. The broad rather flat

strong feature due to N IV A4057 is not seriously

distorted by blends. In addition to the components indicated in the diagram, the Si IV lines

A4088 and X4116 contribute strongly to the blend

at 4100 A.

186
eO 90 5300 l0g0 30 4OSO6O7OaO9O 54O0 020 30 40508070 80 90 55(X) K) 20 A

00 n aO 00 6200 D 20 30 4O 50 0O7D 80 90 6300 O a03O4O5OaOTO8O90 6400 A Figure 2, The spectrum profile of HD191765,

WN6, between 5300 and 6400 A.

40 50 60 TO ao 90 3600 iog0 30.«0 5peo7peo80 3700 y 20 30 «o so S) iO

00 4000 D 20 3O4O3OflO7Oa000 4X30 t>i0 30 40SOaOX>aoeo 4200 C 20 30 A Figure 3, The spectrum profile of HD191765,

WN6, between 3540 and 4200 A.

187
-70 -60 -50 -40 -30 -20 -10 0 10 20 30 40 50 60 70 A (—I

1 1 1 1

1 \ 1 \ 1 1 1 1 1 1—

-70 - 60 - 50 - 40 -50 - 20 -10 0 lO 20 30 40 50 60 70 A Figure 4. Some He II profiles in the spectrum

of HD191765 fitted by a Gaussian profile.

(ii) Most of the emission lines in Wolf-Rayet

spectra have a rounded shape that can be approximated rather well by a Gaussian function. In many

cases the apparent "line" is a blend of several com- ponents. In the case of the He II lines however

the blending is not severe. Some typical He II

profiles in the spectrum of HD191765 can be seen in

Figures 2 and 3. Some of the He II profiles are shown fitted by a Gaussian profile in Figure 4.

The total width at half-intensity of the He II lines

of HD191765 in velocity units remains essentially

constant at about 2 3 80 km/sec for the whole spectrum from X3203 to A6683. This width is the same

for lines of the n = 3, 4 and 5 series. This

rounded shape is the usual case for emission lines

in Wolf-Rayet spectra. The flat- topped lines form

an exception.

188
(iii) The emission lines of the WTH stars are much sharper than those of most Wolf-Rayet stars,

resembling in this respect the emission lines of Of

stars more closely than those of an average WolfRayet star. This point is illustrated by Figures

5 and 6, which show part of the spectrum profile of

three WN7 stars and of the WN6 star HD191765, as well

as by the photographs of spectra presented by Hiltner

and Schild (1966) and by Smith (1968) . Each of the

stars HD9 3131 and HD92 740 seems to be accompanied by

a companion, for absorption cores are visible at Hy

and H6 , although this point has not been noticed by

those doing spectral classification.

Hiltner and Schild have noticed that the emission lines in WN binaries are usually sharper than

those of WN stars not known to have a companion. A

point to be considered when interpreting Wolf-Rayet

spectra, and in particular those classified as WN7

and WN8 , is whether some of the emission lines are formed in gaseous streams or in an extended envelope

around the complete (binary) system. A similar problem is posed by the spectra of Wolf-Rayet stars

90 4600 10 ' » ' 30 ' 40 ' SO ftO ' TO ' »0 ' 80 fKO « ' » ' »T

Figure 5. A comparison of the spectrum profiles

of three WN7 stars with that of a WN6 star in the

spectral region 4580 to 4740 A.

189
HD9313I 1 WN 7

1 ^^/^

HD92740 / I, .J\f WN7 /

00 4>00 Figure 6. A comparison of the spectrum profiles of three WN7 stars wifh that of a WN7 star

in the region 4000 to 4140 A.

which are the nuclei of planetary nebulae. For purposes of interpretation, one must separate the lines

formed in the nebula from those formed in the stellar

atmosphere.

2. Shortward-displaced absorption lines from normal levels.

Shortward-displaced absorption lines are most

simply interpreted as being due to absorption in

that part of an expanding shell which is projected

against the stellar disk. Some lines are known to

be strengthened in absorption owing to dilution effects in an extended atmosphere. Dilution effects

occur because there are metas table levels in the

term scheme of the ion. Such lines are the most

conspicuous absorption lines in Wolf-Rayet spectra.

However a few lines from normal levels are also seen as shortward-displaced absorption components.

The most conspicuous lines are C IV X5801 and A5812

in WC stars (see Figure 1) and N V A4603 and X4620

in WN stars (see Figure 5) . Occasionally in WN

stars He II A4686 and A3203 and the strongest lines

190
of the Pickering series appear as shortwarddisplaced absorption components. Each of these

absorption lines originates from a "terminal" ion;

either the ion is the last possible ion which has

an absorption line spectrum, or a great amount of

energy is needed to reach the next stage of ionization. The ionization potential of the fourth ion

of carbon is 391.986 volts, while that of the fifth

ion of nitrogen is 551.925 volts. Hence the absorption spectra of He II, C IV and N V persist

through a wider range in temperature than the absorption spectra from lower ions. Shortwarddisplaced absorption components are not seen accom- panying all lines in Wolf-Rayet spectra.

Since the absorption lines C IV A5801, X5812,

N V X4603, X4620 and He II X4686, A5411, etc. are subordinate lines from normal levels, qualitative

reasoning suggests that the electron temperature in

the part of the atmosphere where these lines are formed varies directly as the excitation potential

(see Table 1) of the lower level. The 0 VI lines

have not yet been observed in absorption in WolfRayet spectra to my knowledge. The C IV lines are observed in absorption only in WC stars. These

arguments suggest that the electron temperature in

the expanding shell of a WC star is lower than

in the expanding shell of a WN star.

5. Long series of lines from hydrogen-like spectra

According to the arguments of Inglis and Tel'

ler (1939)

,

if the net charge on the nucleus or

core of the ion is Z, the electron density is re- lated to the last visible line of a Rydberg serie:

as follows

log N =

^ e 23.26 - 7.5 log n + 4. 5 log Z.

TABLE 1

Spectrum Level E. P. Lines

C IV 3^S 37.54 volts 5801, 5812

He II n=4 51.00 5411, 4541

N V 3^S 56.54 4603, 4620

0 VI 3^S 79.33 3811, 3834

191
Thus if the hydrogen series breaks off at n = 16,

the He II series will break off at n = 24 and the

C IV series at n = 37. Long series of He II and of

C IV lines are seen in Wolf-Rayet spectra, but this

does not mean that the electron density is particularly low. If the hydrogen series breaks off at

n = 16, the electron density is about 10 The

available observations do not permit a precise es- timate of the electron density in this way because

it is impossible to distinguish the Rydberg series

of He II and C IV to sufficiently high values of n.

4. The composition of the spectrum

It is well known that lines from the ions of

carbon and oxygen dominate the spectrum of WC stars,

whereas lines of the nitrogen ions dominate the

spectra of WN stars. Both types of Wolf-Rayet stars

contain lines from He I and He II , while the hydrogen lines seem to be significant only in the stars

of types WC8 and WC9. Since the lines of the Bal- mer series of hydrogen blend with lines of the

Pickering series of He II , it is difficult to separate the contribution of H from that of He II. If

the decrement from Ha to A5411 to H3 to X4541 to Hy

is fairly smooth, one may assume that the contribution due to hydrogen is small.

Careful inspection of moderate dispersion spectrograms of Wolf-Rayet stars has shown that weak

lines of the nitrogen ions are present in WC spectra, while lines of 0 V and of 0 IV are quite prominent in WN spectra. The C IV blend AX5801, 5812

is definitely present in WN stars, though this blend

is much weaker than it is in WC stars. The corre- sponding multiplet of O VI at 3811 and 3834 A appears

in emission in a few Wolf-Rayet spectra; some of

these stars have been classified as WC, others as WN. Three stars showing these lines are HD115473

(WC5) , Y2 Velorum (WC8 + 07) , and HD104994 (WN3)

There is no sound reason for thinking that the

observed relative intensities of the emission lines

in Wolf-Rayet spectra indicate gross abundance differences between the various spectral classes.

Quite clearly a detailed theory of spectrum formation must be developed before any conclusions about

abundances can be drawn. The particular excitation

processes which are active in the atmospheres must

be considered in detail.

192
III. SPECTROSCOPIC DIAGNOSTICS AND

INTERPRETATION

At present very little quantitative data about

Wolf-Rayet spectra is available for analysis and no sophisticated and fully satisfactory theories of

spectrum formation exist to turn this data into

quantitative parameters such as electron temperatures, electron densities, abundances and state of

motion of the atmosphere. In this section some qualitative arguments will be reviewed which indicate with what temperature and pressure range we

seem to be concerned, and a brief summary will be

given of the physical representations which have

been considered. The geometric properties of the

atmosphere (shape of the atmosphere and its state

of motion) also have an important influence on the

observed spectrum.

a. Evidence Concerning the Temperatures of

Wolf-Rayet Stars

i. From the continuous spectrum

The energy distribution in the continuous spectrum between 3300 A and 10000 A, when corrected for

interstellar reddening, is rather similar to that of

early B-stars or 0-stars. Photometric arguments

similar to those used for OB-stars, the observed

magnitudes having been corrected for the emission

bands, suggest that the color temperatures of WC

stars are near 2x10** °K while those of WN stars are near 3.8xl0'*°K (Pyper 1966). When the color temperature is greater than 2x10** °K, the energy distribution between about 4000 A and 7000 A is not

an accurate indicator of temperature; the color

temperatures might be 10'*°K higher, certainly not

much lower.

The ultraviolet excesses of the WN stars and

the possible infrared excess cannot be explained

by the usual theories of color temperature. Bless,

Fischel and Stecher (1968) , and Wallerstein (1968)

who expands an idea due to Nariai (196 7) , have made

the interesting suggestion that some Wolf-Rayet

stars may be surrounded by a hot corona generated

by the expanding atmosphere colliding with the surroxinding nebula or with the interstellar medium.

It is postulated that x rays are generated in this

corona, and evidence has been compiled that some degree of correlation exists between Wolf-Rayet stars

193
and x-ray sources. The x rays will be accompanied

by high-speed electrons, and it seems possible that

the observed ultraviolet excesses of the WN stars

may be a result of the interactions between radiation and matter in the postulated corona. Since

the spectra of T Tauri stars also have some characteristics similar to the spectra of Wolf-Rayet stars

(ultraviolet excesses, broad emission lines typical

of temperatures near 10'*°K, and a few shortwarddisplaced absorption cores) , one may infer that in

these cases, too, we are observing the effects of

fast moving material impinging on an interstellar

medium or on a shell of not too low density. A

detailed physical picture of what may occur must

still be developed and tested.

The ejection velocity of the shell from the

Wolf-Rayet, Of or T Tauri star cannot be wholly re- sponsible for the ultraviolet excess, since WC stars

have as large velocities of ejection as WN stars,

and both are greater than those of T Tauri stars;

yet only the T Tauri stars and WN stars are observed to have definite ultraviolet excesses. In

the case of WN stars the ultraviolet excesses can- not be attributed to emission in the Balmer continuum. Bless, Fischel and Stecher have suggested that

the expanding atmospheres around Of stars might also

generate hot coronae and x rays. However, Of stars

are not known to possess significant ultraviolet

excesses, although some of them appear to possess

infrared excesses.

2, From the line speotnon

In Section II, b, 2, it was noted that since

shortward-displaced absorption components of C IV

XA5801,12 are seen in WC spectra but not in WN spectra, whereas shortward-displaced absorption components of N V AA4603,20 are seen in WN spectra but not

in WC spectra, the electron temperature in the ex- panding atmospheres of WC stars is lower than that

in the expanding atmospheres of WN stars. In addition somewhat qualitative arguments have been given

by Underhill (1958) , indicating that the observed

difference in composition of WC spectra from WN

spectra may be interpreted as evidence for a different effective temperature of the underlying star,

the effective temperatures of WC stars being about

3x10 '*°K while those of WN stars may be as high as

5xlO'*°K.

The great complexity of WC spectra is largely

194
due to the presence of the strong, rich spectra of

C III, O III and O IV. On the whole WN spectra con- tain fewer lines, and if the well-known N III multiplet AA4634,41 is set aside as being excited

chiefly in an outer nebula, the remaining strong

lines of N IV, N V, O IV and O V give an impression

that the level of excitation in WN atmospheres is

higher than that in WC atmospheres

Using the idea that the spectra are collisionally excited, Bappu (1968) has estimated excitation

temperatures for five Wolf-Rayet stars. His results

are given in Table 2.

The electron temperatures in the atmospheres

of Wolf-Rayet stars are not known well. The most

probable range seems to be 2.4xlo'*°K to about

5xl0''°K in the part of the atmosphere where the

rounded emission lines are formed. Estimates made

by others can be found in the references given in

Section III, o, 2.

b. Evidence Concerning the Electron Density in

Wolf-Rayet Atmospheres

In Section II, h, 3 it was pointed out that

the Inglis-Teller formula allowed one to place a loose upper limit of 10 on the electron density.

No sharp forbidden lines , such as are observed for

planetary nebulae, appear in the spectra of WolfRayet stars. Furtheinnore one does not observe

sharp, narrow absorption lines. The electron density is probably well above that in planetary nebulae. Model atmosphere calculations for 0-and Bstars suggest that electron densities in the neighborhood of 10^^ to 10^^ may be appropriate for the

inner parts of Wolf-Rayet atmospheres. Wallerstein

TABLE 2 HD192163 WN6 32 400°K from lines of He II

HD165763 WC6 56 900 from lines of C IV

HD192641 WC7 + Be 51 600 from lines of C IV

HD192103 WC7 38 000 from lines of C IV

HD184738 WC8 24 000 from lines of C IV

195
(1968) estimates 10^^ from the x-ray intensity of

a source which he relates to HD211853.

c. The Types of Physical Process to he Considered

Three aspects of the interactions between

radiation, atoms and ions that lead to the phenomenon we call a Wolf-Rayet spectrum must be considered. The first is the mechanism of spectral line

broadening, while the second is the radiative transfer process itself, and the third is the geometry of

the situation.

I . Physical mechanisms of line broadening

Three major processes may be considered: Stark

effect, Doppler broadening, and the effects of non- coherent electron scattering. Because the shapes

of the He II lines appear to be independent of the

lower quantum number (n = 3, 4 or 5) and of the

upper quantum number (n = 4 to 14 or so) , one may

conclude that Stark broadening is not an important

mechanism. If the electron densities lie in the

range 10^^ to 10^^ as surmised, significant Stark

broadening would not be expected. The fact that

the shapes of unblended lines in any one spectrum

have the same halfwidth and shape, on a velocity

scale, and the fact that most of the lines have a rounded, Gaussian shape make it seem probable that

Doppler broadening is an important cause of the

observed line shapes. If large motions exist in a Wolf-Rayet atmosphere, say a Gaussian distribution

of velocities with a most probable velocity of the

order of 1000 km/sec, the atmosphere would be optically thin,, and profiles of the observed shape

might result. However the broadening resulting

from electron scattering will also produce a rather

rounded profile from an originally fairly sharp emission line. One critical point to be determined is

how great an opacity of electrons would be required

to produce the observed line shapes from an initially

sharp line. It is true that the emission lines are usually rather wide in Wolf-Rayet stars (this width

varies considerably from star to star) , but the absorption lines. He I X3888 in particular, are not

very wide. One can perhaps limit the amount of electron scattering which is permitted by making sure that the absorption lines remain as narrow as they

are observed. This problem of the balance between

Doppler broadening and broadening by electron scat196
tering cannot be solved without considering the

geometrical figures of the atmosphere.

2. The radiative transfer processes

The first attempts to analyze Wolf-Rayet spectra in order to find temperatures started from the

extreme hypotheses of nebular conditions (Beals

1934, 1940) and of thermodynamic equilibriiam (Aller

1943) , with only radiative processes considered to

be of importance. The results were unsatisfactory

and contradictory, and it was evident that more sophisticated theories would have to be developed.

Attempts were made to improve the purely radiative

theories by Zanstra and Weenen (1950) , Miyamoto

(1952) and by Rublev (1964) . It became evident that

the spectrum formation process in Wolf-Rayet atmospheres was more closely akin to that in Be stars

or shell stars than to that in planetary nebulae.

In the purely radiative theories it is usually

assumed that the quantity and quality of the ex- citing and ionizing radiation at wavelengths below

the Lyman limit can be described by a single temperature and a blackbody distribution. A consistent solution is sought to account for the observed

equivalent widths of a few emission lines in terms

of an effective temperature and a set of abundances.

Little or no attention is paid to the shapes of the

lines

The importance of collisional excitation processes was emphasized by Thomas (1949) . Quantitative application of these ideas has been done by

Bappu (1958, 1968). The problem has been looked

at in some detail by Code and Bless (1964) who have

shown that protons and alpha particles moving with

velocities of the order of 1000 km/sec have sufficient energy to generate the observed excitation.

Collisional excitation is an attractive way of obtaining the observed wide range of excitation which

appears in nearly every Wplf-Rayet spectrum.

It is difficult to interpret Wolf-Rayet spectra accurately and in a quantitative manner, for one must separate the spectrum into groups of lines which

are formed under the same atmospheric conditions.

Many years ago Cecilia Payne-Gaposhkin (1935) drew

attention to the fact that Wolf-Rayet spectra are a collection of spectra from different sources. How

to effect the needed separation is not yet certain.

Great practical difficulties exist because there are,

indeed, very few lines which are not seriously

blended, and very little quantitative spectropho197
tometry has been done using spectra of adequate

spectral purity.

Z. The geometvioal aspeots

The concept of a simple, spherical, low density expanding atmosphere (Deals 1930, Chandrasekhar

19 34) seems to be adequate for interpreting the

shapes of a few lines. However as 0. C. Wilson

(1942) demonstrated by considering binary stars,

the expanding-shell hypothesis is inadequate for

explaining most of the lines in Wolf-Rayet spectra.

One must consider that the rounded emission lines

are formed in a rather compact atmosphere near the

stellar photosphere. The effects on the line profiles of various geometric shapes and of various

velocity fields have been explored by Bappu and

Menzel (1954); Rublev (1960, 1962, 1963) has used

Sobolov's treatment of the moving atmosphere problem to explore further. The fact that some shortward-displaced absorption components and some flat-topped emission lines exist seems to be definite evidence that matter is streaming from WolfRayet stars. This streaming appears to be steady;

the expansion velocity in any one star does not

change. Velocities have been observed from a few

hundred km/sec to about 2800 km/sec. Radiation

pressure in the far-ultraviolet resonance lines may

be the driving force. This question has not been

thoroughly explored.

IV. A MODEL OF A TYPICAL WOLF-RAYET ATMOSPHERE

The spectroscopic detail observed in apparently

single Wolf-Rayet stars such as HD192103 (WC7)

,

HD191765 and HD192163 (WN6) , can be understood, at. least qualitatively, in terms of the following model

which consists of three parts.

(i) First there is a photosphere which ra- diates in the continuous spectrum like an 0-star.

For a WC7 star the effective temperature may be

about 3x10 ''°K, while for WN6 stars the effective

temperature is probably at least 4x10 '^^K. A separate source of radiation, perhaps a hot corona, is

required to generate the ultraviolet excess of the

WN stars. Since most Wolf-Rayet stars are about as bright visually as 0-stars (i.e., -4 < My 5 -5), the

photosphere must have a radius like that of an 0-

198
star. Seven to ten solar radii seems to be a rea- sonable range.

(ii) Second there is an inner compact atmosphere where all the emission lines with rounded

profiles are formed. This atmosphere is rather

opaque in line frequencies but transparent in con- tinuum frequencies between 3000 and 10000 A. If

the width of the lines is chiefly due to Doppler

broadening , the ions in this part of the atmosphere

have motions such that <v> is of the order of 1000

km/sec. The thickness of the atmosphere may be one

to two solar radii. The particle density probably

lies in the range 10^^ to 10^** particles per cm^

(Hi) Outermost there is an expanding lowdensity atmosphere in which dilution effects and

monochromatic fluorescence effects occur. This

atmosphere can be considered to be due to the boiling off of the inner atmosphere. The evaporation

could be generated by the Maxwell tail of the ve- locity distribution in the inner atmosphere. A

typical radius is 17 solar radii (cf . , the dimensions estimated for the eclipsing variable V444

Cygni) . In WC stars the level of excitation in the

expanding shell is 30 to 40 volts; in WN stars it

is 50 to 60 volts.

The WN7 and WN8 stars appear to be very luminous. My -6.5, like the OB-supergiants and P Cygni.

Insufficient information exists to sketch a useful

model.

V. CLUES REGARDING THE EVOLUTIONARY

STAGE OF WOLF-RAYET STARS

The visual absolute magnitudes of most WolfRayet stars lie between -4 and -5. Only the WN7

and WN8 stars appear to be as bright as -6.5. If

the bolometric corrections are about the same as for 0-stars, about 3 magnitudes, the Wolf-Rayet

stars in general are about as bright bolometrically

as 0-stars, while the WN7 and WN8 stars are similar

to OB-supergiants. From the seven double-lined

spectroscopic binaries for which a mass ratio has

been estimated, it is quite clear that the mass of

the Wolf-Rayet star is about 1/3 that of its early

B- or 0-type companion. In no case is the mass greater than that of the 0-or B-star, or even equal

to it. This means that the mass of a Wolf-Rayet star

probably falls in the range 4 to 10 solar masses. Nothing is known directly about the masses of WN7

and WN8 stars.

199
Arguments concerning the distribution of WolfRayet stars put their ages at 10® years or less.

Arguments concerning the sizes of the interstellar

ring nebulae observed around some stars place their

ages in the range 10^ to 10^ years. These discussions suggest that WC stars are older than WN stars.

The characteristic spectra of Wolf-Rayet stars

are a direct result of the interactions between

radiation and matter in plasmas having particle densities of the order of 10^^ to 10 electron temperatures of the order of 3xl0''°K to Sxlo'^^K, and a random velocity field with <v> of the order of

1000 km/sec, as well as a low-density gas streaming

outward with a velocity between about 5 00 and 2 800

km/sec. The temperature and velocity vary somewhat

from star to star.

Since the normal Wolf-Rayet spectra have many

details in common with both T Tauri stars, which

are generally acknowledged to be in the first stages

of contraction, and with the central stars of planetary nebulae, which are generally acknowledged to be

at an advanced stage of evolution, it is clear that

the characteristic appearance of Wolf-Rayet spectra

is not a sufficient criterion to establish the evo- lutionary stage of these stars. The typical emissionline spectrum is merely a statement that the physical

conditions described in the preceding paragraph exist

in a plasma around the star. This set of conditions,

apparently, can come into being at more than one stage

in the lifetime of the star. What process makes just

those physical conditions appear in the outer layers

of a star with the result that the spectrum is of

"Wolf-Rayet type" is at present unknown.

The above arguments concerning the ages of

Wolf-Rayet stars and the fact that the Wolf-Rayet

star is always the least massive star of the binary

system (in those cases when two spectra are observed) make it seem probable that the Wolf-Rayet

phase occurs while the star is contracting to the

main sequence. On the other hand, it is possible

that the mass ratio near 1/3 is the result of mass exchange in a close binary system. Then one must

conclude that the Wolf-Rayet phase corresponds to

a stage fairly far along in the evolution of a star. Those Wolf-Rayet stars which are the nuclei

of planetary nebulae must be considered to be at a late stage of evolution. These stars seem to be

different from "normal" Wolf-Rayet stars. The

meager evidence about their visual absolute magnitudes places My near 0 or +1 , and their radii appear to be of the order of the solar radius.

200
If one excludes the central stars of planetary

nebulae from the discussion, then two choices are open for the evolutionary stage of Wolf-Rayet stars:

either (1) they are still contracting to the main

sequence, or (2) they are in a post-main sequence

stage of development following mass exchange in a binary system. There appears to be no sure way of

deciding between these alternatives on the basis

of spectrum alone. It is true that the observed

visual absolute magnitudes imply higher bolometric

absolute magnitudes than the theory of contracting

stars of known mass predicts. However this con- clusion is true only if the assumed bolometric cor- rection of about 3 magnitudes is correct. It is

necessary to determine the bolometric corrections

for Wolf-Rayet stars.

REFERENCES

Aller, L. H. 1943, Ap , J., 97,^135.

Bappu, M. K. V. 1958, Etoiles a Raies d'Emission

(Universite de Liege), p. 40.

Bappu, M. K. V. 1968, Q.J,R.A,S., 9, 83.

Bappu, M. K. v., and Menzel , D. H. 1954, Ap . J,^

119, 508.

Beals, C. S. 1930, P.Z?.^.0.^ _4 , 294.

Beals, C. S. 1934, P.D.A.O,, 6, 95.

Beals, C. S. 1940, J.R,A,S,C., 34., 169.

Bless, R. C. , Fischel, D. , and Stecher, T. P. 1968,

Ap, J. (Letters), 151, L117.

Chandrasekhar , S. 1934, M.N,R.A.S, 9±, 522.

Code, A. D. , and Bless, R. C. 1964, Ap . J., 139 , 787.

Hiltner, W. A., and Schild, R. E. 1966, Ap . J.,

143, 770.

Inglis, D. R. , and Teller, E. 1939, Ap . J., 90, 439.

Kuhi, L. V. 1968, Ap . J,, 152, 89.

Miyamoto, S. 1952, Cont. Inst. Ap , Kyoto, No. 24.

Nariai, K. 1967, Pub. Astr. Soo. Japan, 19., 564.

Payne-Gaposchkin, C. H. 135, Trans. I.A.U., 5, 187.

Pyper, D. M. 1966, Ap . J., 144, 13.

Rublev, S. V. 1960, Astr. Zh. , 37., 828; Soviet

Astr.-AJ, 4, 780.

Rublev, S. V. 1962, Astr. Zh. , 39_, 879 ; Soviet

Astr.-AJ, 6, 686.

Rublev, S. V. 1963, Astr. Zh. , 40., 643; Soviet

Astr.-AJ, 7, 492.

201
Rublev, S. V. 1964, Astr. Zh. , 41, 63, 223; Soviet

Astr,-AJ, 8, 45.

Smith, L. F. 1968, M.N,E.A.S.^ 138, 109.

Thomas, R. N. 1949, Ap . J., 109, 500.

Underhill, A. B. 1958 , Etoiles a Raies d ' Emission

(Universite de Liege), p. 26.

Underhill, A. B. 1968, Ann. Rev. Astr. and Ap . ^ 6_, Wallerstein, G. 1968, Ap. J. (Letters), 151, L121.

Wilson, O. C. 1942 , Ap . J., 95., 402.

Zanstra, H. , and Weenen, J. 1950, B.A.'N., 11, 165.

202
DISCUSSION

Chairman: Lawrence H, Alter

[Ed. comment: The original plan of the symposiiim was that Part B should contain all the empirical material on spectral features and that Part C

should be a digest of ideas on (1) the atmospheric

conditions required to produce such spectra and (2

)

the models and physical effects required to produce

the atmospheric conditions. Unfortunately, the

literature contains few detailed quantitative investigations of the kind required for Part C. Miss

Underhill therefore thought it best to explore the

observational consequences of various suggestions

for the interpretation of specific spectral features in terms of isolated atmospheric conditions.

Hence much of her discussion supplemented Kuhi '

in presenting further observational material illustrating her points.]

Alter: We shall start by considering any additional observational material relating to the problem of interpretation. Then we can continue to discuss the interpretation of the spectrum, and finally

we can consider the models.

As the first item under "observational material" , I would like to mention an observation by

Gordon Wares and myself of the object n Carinae,

which bears on the effect of electron scattering in

an extended envelope. This object consists of a brilliantly red nucleus (which has a soft appearance under the best seeing conditions) 2 to 2.5" in

diameter and a surrounding shell of about 10" in diameter. The spectrum of the stellar nucleus has been

studied by many observers (see, e.g., Thackeray 1953,

M.N,R.A.S.^ 113 , 211; Aller and Dunham 1966, Ap . J,^

146 , 126; Rogers and Searle 1967, M.N,R.A.S. ^ 135,

90) . It consists of a background continuum, the re- markable properties of which have been studied by

Rogers and Searle; broad emission lines largely of

hydrogen and ionized iron; sharp emission lines

superposed on the broad emission lines; and occa- sional absorption lines on the violet edges of the

emission lines. There is a number of forbidden

lines of Fe+, Ne+ and S+ , but none of 0+. The broad

emission lines and violet absorption edges resemble

the spectra of Wolf-Rayet stars, but the latter show

no counterpart of the sharp emission lines.

203
It is generally stated in the literature that

the spectriim of the shell resembles that of the

nucleus, so we decided to attempt a quantitative

comparison. Figure 7 shows a series of spectra

secured with a dispersion of 39 A/mm with the 60-

inch reflector at Cerro Tololo. The longest exposure on the nucleus is about 7 minutes and that on

the shell (which is flanked by the companion spectrum) is about 33 minutes. This plate covers the

region from 5015 to 4070 A. Additional spectrograms cover the region 4100 to 3200 A and 5000 to

6 700 A. The light from the shell has been regarded

as nuclear light scattered by electrons, and we feel

that this explanation must be the correct one. The

emission lines in the shell spectrum all appear to

be derived from the central nucleus. The sharp lines

of the core are gone; nothing is left but diffuse

features. The absorption lines, on the other hand,

appear to be strengthened in the shell, which has

thus Wolf-Rayetized the spectrum of the n Carinae

nucleus. We appear to have here a graphic demonstration of the importance of electron scattering,

which has been suggested by Miinch, by Kopal and

Shapley, and by Thomas as a significant mechanism

in Wolf-Rayet stars. A quantitative study of these

spectrograms is in progress.

Smith: You've shown that the violet absorption

edges are more intense when you look toward the

shell than when you look straight at the central

star. But we've been interpreting these absorption

edges as coming from the part of the shell between

us and the star; so why should they be stronger when

you're not looking at the star?

Alter: They are stronger only relative to the

emission in the adjacent continuum. In an absolute

sense, in terms of equivalent widths, they are weaker. This of course involves a radiative trans1 1 1 Hi r^T^WT T I

11 II ill Hill

II W III Hill

IllHi i II : ^ 1

lllll'l i II : 1

1 II 1

1 1 1 1 ^

Figure 7. A comparison of spectra of the central nucleus of n Carinae (top) and of the outer

shell (bottom) , taken with the 60-inch telescope

at Cerro Tololo. Wavelength increases to the right,

204
fer problem; one has to interpret line shapes as well as intensities. We're trying to do the photometry on these emission and absorption profiles now,

in order to compare the two effects. I think the

detailed quantitative interpretation will require

the solution of a transfer problem.

Thomas: If we buy the conventional picture of

an absorption line formed only in the material between the stellar disk and the observer, then an absorption edge will appear in the shell spectrum

only if it has been scattered through 90°. What

you are saying is that the effect of electron scattering is greater on the continuum and on the emission line than on the absorption edge - the absorption is washed out less than the emission.

Alter: That- would be an obvious interpretation, but one also has to worry about excitation of

lines in the outer shell. The reason we don't think

there is much direct excitation of Fe II is the ab- sence of lines with sharp cores in the shell spectrum.

Smith: Are you telling me that the conventional picture with the absorption core coming just from

that part between us and the star is an utterly un- realistic simplification? We must include the scattered light from the shell?

Underhill: It's not utterly unrealistic. But

in this case there seems to be a rather dense shell

which is optically thick in strong lines such as hydrogen; hence they will be self-absorbed. If you

look at very strong lines in an iron arc spectrum,

you'll find self-absorbed cores, the shape of which

depends upon the motion. Here we have exactly the same thing. The shell is optically thick - very

thick - in the hydrogen lines; you cannot assume it

has the same thickness in all lines. It's very

simple

Payne-Gaposohkin : Some facts on n Carinae

should be mentioned in connection with the production of these line profiles. The first observation

of n Carinae, made in 1887, by Agnes Clark, showed

bright lines, whereas a spectrum taken at Harvard

in 1889 showed only absorption lines. It resembled

an F-supergiant with rather hazy lines. Some of

the lines, such as those of hydrogen, showed faint

bright edges.

Alter: The surrounding shell is usually interpreted as the result of an outburst in 1843 when the

star reached maximum brightness. Unfortunately we have no spectroscopic observations so we really

don't know what happened. The outburst in 1895 was

205
less violent. The object has apparently settled

down now to the spectrum I described. I have very

carefully referred to this as an object rather than

a star because of its large diameter; what it is, we don't know. If the density in the shell were

higher, it might have resembled a WR star.

Wvubel: One type of datum that is obviously

missing and that might contribute to our knowledge

of electron scattering is measurements of the polarization of the emission lines. Kuhi has already

said this is beyond his techniques at present. Does

anyone else have any information?

Johnson: Hiltner obtained polarization data

on V4 44 Cyg some years ago, but I don't know of

any work since then.

Kuhi: Hiltner 's work was all in the continuum. What about the lines?

Johnson: Apparently there is no such work.

Polarization measurements can be made quite accu- rately now, so someone should certainly go into this

problem. It might settle the question of electron

scattering in these stars.

Allev: Possibly we could now say something

about radio and x-ray observations of WR stars.

Fecker-Wimel : At the symposium on planetary

nebulae last year, Davies announced that two normal

WR stars had been observed at 11 cm. This should

be very interesting to us because it should corre- spond at least to conditions in the solar chromosphere . Ej ellming : Was the resolution sufficiently

good to be sure the radiation was not coming from

a nebular shell?

Peckev-Wimel : They were two normal WR stars

without nebulosity.

Alter: This could certainly be very significant because of the associated suggestion that WR

stars are x-ray sources. We have already discussed

this in the last session with respect to the two

southern objects with strong O VI emission lines.

Steoher: This is an important question which

might tell us something definitive about WR stars.

The thermal expansion we have been talking about

corresponds to 10 ^°K. Eventually the particles will

run into the interstellar medium and should produce

X rays. There is also the possibility of line emission in the x-ray region, which I think we are just

able to tackle instrumentally

Thomas: Let's be very clear what you are talking about here. When you speak of a thermal ex- pansion at 10 ^°K, I assume you are referring to my

206
figure of 20 years ago, when I said that if_ you want

the WR atmosphere to be maintained in hydrostatic

equilibrium with a temperature, it has to be 10 ^°K.

But I also said I didn't believe this temperature;

I preferred instead a non-static atmosphere supported

by some sort of jets. But this points up our real

problem in trying to infer atmospheric structure

from excitation levels. In the visual spectrum of

the Sun, for example, there is nothing to indicate

temperatures of the order of 10^ °K in the corona. And it is hard to infer from thermal effects at

10^ °K the existence of x rays and cosmic rays, such

as are observed to come from the Sun. So for the

Sun, you have a choice of looking for thermal emission of X rays in localized "very hot spots", or of

looking for some non-thermal-emission processes of

the sort that Biermann and Liist have been seeking

for many years. The same problem arises in WR stars,

and this is why I have been trying to pin you down

on excitation levels in various parts of the spectrum as a reflection of excitation levels in various parts of the atmosphere. Anne Underhill says

she doesn't want to go higher than 5xlo'*°K in the

continuum. In the visual spectral regions of the

line spectrum, we see N V and possibly 0 VI . I don't recall Stecher pointing out higher ionization than this in the rocket UV for X > 1150 A. So

if I assume collisional ionization in the usual nonLTE approach, the line excitation corresponds to some 2 to 3x10^ °K. But of course if Tg can be 10^ °K

in the Sun, it is hard .to imagine that with the

evidence of an even greater supply of mechanical

energy in WR stars we don't get Te at least this

high. The question is, how much higher? There has

been much speculation on the kind of equilibrium

configurations that can be reached as a function of

mechanical energy input and atmospheric opacity,

but I don't think any definitive conclusions have

been reached. So if you want to start talking about

x-ray sources, we are left with the same problem we have for the Sun. We will have to investigate atmospheric models as a function of mechanical energy

input and of the radiative input, which fixes the

photospheric model. Until this is done, we have no basis for distinguishing between thermal and nonthermal explanations of what we observe.

Steoher: The Doppler shift for, say, C IV

corresponds to 0.25 meV C IV particles.

Thomas: What counts for excitation is differential motions. Don't forget that you won't get

0.25 meV collisional excitation from 0.25 meV heavy

207
particles. It's the old story of heavy particle

excitation not being as efficient as electron excitation . Feckev-Wimel : We don't need to be so erudite.

The WR stars with x-ray spectra have O VI; if they

have O VI , that means they can have C V and you can get plenty of x rays from C V.

Smith: I would like to clarify this question

of the presence of O VI lines in WR spectra. It

has been said here that 0 VI is only observed in

the spectra of a very few stars. However, in the

Hiltner-Schild Atlas (Figure 24) [Figures 24-28

appear at the end of Part B, pp. 175-179] , you can clearly see a pair of emission lines at XX3811 and

3834 which correspond to the 0 VI doublet. You can

see that they are present in the WC6 and WC7 spectra; in my spectra of WC5 stars they are also present.

Undevhill: There are also possible blends of

0 III and 0 IV. I remember studying HD192103, and

there were many lines in that region. I could not

convince myself that the 0 VI lines were definitely

present among the other lines, most of which were

due to the less-ionized 0-ions. I agree with Lindsey that in some of the hotter stars, O VI looks more definite.

Smith: HD192103 is a WC8 star, and I am not

suggesting the presence of the 0 VI doublet in that

star. I went through Edlen's list {Vistas in Astronomy ^ 2, 1456) and the revised multiplet tables,

and I concluded that this XA3811, 3834 pair is a blend of 0 VI and He II. Since the other emission

lines in the same He II series on either side are very weak, the observed emission must be dominantly

0 VI. The appearance of the 0 VI doublet is similar

in the WC5 , 6 and 7 subclasses, becoming stronger as you go from WC7 to WC5 and are stronger still in the

planetary nuclei mentioned by Aller.

Johnson: 1 would like to re-emphasize Aller 's point that the spectra of the two proposed x-ray

sources, NGC 5189 and GX3+1, are dominated by the 0 VI

doublet and by He II X4686. Since these objects

are associated with x rays, they must either be

very hot or contain a very hot source of radiation.

Furthermore, the lines are very broad; I recall

something like 10000 km/sec being mentioned.

I would like to discuss the problem of classifying a WR star that has neither C nor N lines

visible in the spectrum. Maybe they are not WR

stars, and we have to discuss them as examples of

0 VI . It is my impression that we have never reach208
ed an absolutely rigorous method of dropping stars

into a slot labelled "WR star".

Thomas: 1 disagree. At the outset, I proposed

four criteria that were broad enough to include

all the objects we have just been discussing, yet

specific enough to exclude a number of others on which questions have been raised at various times

during this symposium. Whether you like my four

criteria is another matter; but they do seem to me

to summarize the unique features of the WR category

of objects or spectra.

But let's return to the critical question. In

Anne's summary, she stated that 0 VI was present in

only two stars. Now I gather that if Lindsey is

right, it has been observed in many more. If you

are right, Lindsey, how many stars show O VI?

Smith: Everything hotter than WC7. That means

26 known classical WR stars in the Galaxy and 15 in

the LMC, including binaries in which the 0 VI lines

may not always be visible due to "drowning".

S child: I must disagree on one point, and that concerns the He WR star HD6327 (Figure 28) . It

shows all the features of WR stars except that it

does not show any lines of C, N, or 0 , at least in

the blue and violet regions. We have only one spectrogram covering the region about AX3400-6500, and we see nothing but He.

Thomas: It's not a WR star by my definition;

it may be by yours.

Payne-Gaposohkin : You can vaguely see quite

a number of absorption lines on your spectrum; have

they been measured and identified?

Schild: As I recall, all the absorption lines

are interstellar.

Johnson: It is interesting to me that no one

is concerned about the absence of ifiterstellar absorption lines in WR spectra. If they aren't present, could it be that many of these stars are closer than we think or that they lie in special

regions of interstellar space?

A. Cox: Given the evidence for very high ex- citation levels in the region of line formation and

for rather high Te in the region of formation of the

continuum, it would be interesting to know what little

you can get from spectra about the hydrogen abundance. Perhaps there is significant mass-exchange

and perhaps it leads to changes in internal abundance that affect the evolutionary track of the

star; any information we could get about the hydrogen

abundance would be useful as a check on these speculations . 209
Underhill : All you can say is that in the

cooler WR stars - certainly in the WC7, 8, and 9 stars - you do get bumps in the Balmer decrement.

So you might possibly say there is a little hydrogen

there. I do not know of sufficient material for the

WN7 or 8 stars to be able to check this. I've tried

to do this in the broad-line WN6 stars, and the decrement was very slow and smooth, so I conclude the

lines are certainly 90% He II. But again this only

says H is there; I don't know how you would estimate

how much.

Thomas : How can you expect to observe H in

such hot objects? How far out in the atmosphere

must you go before things get cooler; to the interstellar medium?

Underhill : That's the problem, because the H

you observe is essentially in the interstellar medium.

A, Cox: Another point. You say that the WN

stars have hotter photospheres than the WC stars.

Does this have anything to do with the hydrogen

abundance in the photosphere; is it a model atmosphere problem?

Underhill: I don't think so, but the model

atmospheres on which I am basing my arguments are essentially hydrogen models. If you make a He model,

and compute the intensity distribution, it's going

to come out about the same, because the opacity still

drops off as V ^ . The He opacities come from the

n=3 and 4 levels of He I and He II , and these are treated as hydrogen-like.

A. Cox: The point is that the ionization equilibria for H and He have a different temperature

dependence. I would think that different relative

amounts of H and He would give you different effective temperatures. Do I gather it is possible to

say that H has an abundance somewhere from 0 to 70%,

but we just don't know what?

Aller: If you admit HD45166 is a WR star - maybe Anne won't - then H is certainly present. It seems to be on the borderline with the Of stars, and

I don't think we can easily reach any conclusion because of its extremely high temperature. But I would expect to find some evidence for H in some of

these objects.

Beals: I have made some measurements of the

Balmer decrement, and I am still convinced that there

is a considerable amount of H in the WR stars.

Stecher: If there is no H, must there not be some other sources of opacity? The models of Hunger

and Klinglesmith show that C is important.

210
Thomas: I am bothered by your dependence on the decrements of lines. It seems to me you are ignoring the transfer problem, hence the problem of

line formation. Do I understand that you are simply

looking at the observed decrement, comparing it to

the He recombination' spectrum you would expect from

a thin atmosphere, and deciding on the presence or absence of H according to whether the two agree?

Underhill : Suppose we do assume a thin atmosphere. Suppose further we have a ratio of Ha to

H3 of between 5 and 3. Then we conclude there is

hydrogen. On the other hand, suppose we have a ratio of between 4 and 2. Then we have two choices:

Either there is very little hydrogen, or there are appreciable self-absorption effects. In other words,

if we have a flat decrement, there is no unique interpretation. But if we have a steep decrement,

I'll gamble and say there is a fair amount of hydrogen.

Thomas: This is indeed about as crude an analysis as it is possible to make. It ignores all the

physics of line formation. If there is any opacity

at all in the atmosphere or any collisional excitation I would not put much trust in such arguments.

Aller: It's quite evident that there is a lot

of self-absorption in these stars, as I emphasized

in my 194 3 paper. You have to compare the Pickering and the Balmer series and the Balmer decrement

as best you can, and also other lines for which

good f-values are available. There is also a further point concerning the infrared. Have Low or others working in that region done anything on WR

stars?

Kuhi : I have suggested to them. that they look

at these objects, but I don't know if they have

found anything yet.

Aller: So the infrared is one of the things

we should be studying with a large telescope. Then

there is the related question of the infrared and

ultraviolet excesses.

Kuhi: The first thing I'll do after this

symposium is apply the new calibration of Hayes to

these excesses. I think we can remove some of the

infrared excess, but it's going to be difficult to remove all the ultraviolet. I refer to the excess with respect to a normal 0-type star. The second

point I want to make is that the comparison between

the WR stars and the supergiants should be carried

a little further. It is possible that whatever

is producing the infrared excess in the WR star is

also producing the same effect in OB-type super211
giants. To the best of my knowledge, no one is

studying this problem.

Smith: Let me see if I understand the current

situation. In the old days, before the correction,

the models agreed with OB-stars and not with WR

stars. Now the models agree with WR stars, but not

with OB-stars.

JJndevhill : No. The models never agreed in

detail with OB-stars. There were always small but

disturbing discrepancies, which started people

worrying about the absolute calibration. This is now being straightened out.

Smith: Does this mean the energy distribution

of a WR star is the same as of an OB-star?

Underhill : Not quite, although it is quite

close. The WN's have a definite UV excess, while

the WC's track quite closely. The point I want to

make about the WR stars is that as recently as 10

years ago, people were using photographic photometry to get temperatures of 1.2xl0'*°K. These values

are now recognized to be greatly in error. Now

suppose you make a standard model atmosphere for

5x10 '^''K. If you include line-blanketing it gets

chopped off below X1500. Stecher's observations

show that in the region XX1500-3000, the model fits

the observations very well. What we have just been

talking about, however, is the region X ^ 3300 A.

In this region the 0- stars and the WC's track quite

well, but the WN's go much too high. It's hard to

say what the small corrections may be at the infrared end of the spectrum, as the blackbody curve is

small here, and small deviations correspond to large

changes in effective temperature.

Steoher: The interstellar-grains people should

be able to help us with the infrared correction.

In my own spectra, I don't see the usual lines of

an 0-star; thus I conclude that I'm looking at a WR star. I think the best value for the temperature is about 4xio'*°K.

Roman: I'm still confused about the UV excess. WR stars appear to resemble supergiants; First, do we have any good comparisons between WR stars and

supergiants? Second, do we have good models for

supergiants? Finally, is the great number of emission lines found in the visible part of the WC spectrum related to the apparent differences between

the WC's and the WN's in the UV?

Vnderhill: The answer to all three questions

is no. To determine the continuum for the WC's is

very difficult, although Kuhi has solved it fairly

well. What convinces me of a real UV excess in WN '

212
is that I can get a WN spectra down to A 3100 even when I cannot get spectra of WC ' s or 0-stars.

Alter: I am puzzled by the alleged connection

between the UV excess in WN stars and in T Tauri

stars

Kuhi : The UV excess in T Tauri occurs down

around the Balmer continuum, increases in intensity

toward the UV, and then drops off. In the brighter

T Tauri stars, the UV excess correlates with the Ha

emission, suggesting a Balmer origin. However,

Walker has some spectra of fainter T Tauri stars

which introduce a serious difficulty: whereas the

Balmer lines decrease in intensity and disappear

long before the Balmer continuum is reached, the UV

continuous emission remains. In T Tauri stars there

is also a sort of blue continuous emission which acts

to fill in the absorption lines and makes them difficult to see. So I don't know if one can really say

the WN and T Tauri color excesses are related.

Bess ell: The color excesses for the T Tauri

stars are determined with respect to the energy distribution of an F- or G-star. So a T Tauri star with

a UV. excess is an object like a G-star with enhanced

UV radiation. This suggests a model for the T Tauri

star of a cool star surrounded by a hot shell, which

is similar to the WR model of a hot star surrounded

by an even hotter shell. However the recent dynamical model of a T Tauri star constructed by Larson at

Cal Tech is just the opposite - a hot star surrounded

by a thick cool envelope. So perhaps we should think

of a blue star with a red deficit, rather than a red

star with a blue excess. Underhill : But in addition to the shell features the T Tauri stars show a very definite underlying absorption line spectrum of G- or F-spectral

type. Perhaps I introduced a red herring in pointing out the similarities between them and WR sta;rs.

My point was simply that both types of objects appear to have an expanding extended shell; they differ in their levels of excitation.

Bessell: But you really introduced T Tauri

stars to support the hypothesis that WR stars are also coming onto the main sequence.

Underhill: I introduced them because I feel

we know nothing about the kind of spectrum produced

by an expanded, highly turbulent atmosphere. I asked myself what I would see if I dropped the ex- citation level in a WR star: the answer is the Ca"*" H and K lines and the H lines; they will be broad,

and they may have displaced absorption cores - hence

the T Tauri stars. I agree that the UV excesses are

213
not a good point of comparison, since that of the T

Tauri stars comes from the Balmer emission which

doesn't exist in the WR stars.

Alter: If there is no further explicit discussion on the continuum, I suggest we move on to

the interpretation of the lines, where we have

three basic points: the relation between intensity

and the population of energy levels; the question of

thick versus thin atmospheric shells, which is con- nected with the problem of self-reversals; and the

problem of radiative transfer in a moving atmosphere,

where you may also have curvature effects.

Thomas: I am disappointed that thus far in the

discussion of the interpretation of line spectra we have heard nothing that wasn't talked about 20 years

ago. Anne's discussion of Bappu's 1968 work on T^x

is simply a rehash of the critique I made in 194 8 of Aller's 1943 work. What we have done in these 20

years is .develop quantitative methods for discussing stellar atmospheres and their spectra in terms

of the actual physical situation and not simply an assumed set of physical relations. In order to es- tablish some guide for our analysis, let me go back

to my introductory remarks of the first day and ask

you to consider what really characterizes the WR

spectrum, makes it unique.

Three things stand out in the WR spectrum: (1)

It is an emission-line spectrum; (2) the lines are very broad; (3) the excitation level of the lines

is much greater than that of the continuum. I don't know of any classical stellar model that predicts such a situation. The closest situation for

which we do have models is the solar chromospherecorona, and these models are highly non-classical.

If I want to produce emission lines, I must introduce a mechanism that either puts more energy into

the lines than the continuum, or removes more en- ergy from the continuum than the lines. The Schuster

mechanism does the latter, but Katharine Gebbie and

I have shown that it will not be very effective

under these conditions. Various kinds of fluorescence can do the former, but whereas Anne re- lies heavily on them, I believe they are too fortuitous to explain the whole WR spectrum. The same objection applies to explanations in terms of

nebular- type recombination spectra, even if such

spectra were not already ruled out by point (3)

Underhill : I don't see how you can make the

last remark when we know nothing about the continuum. Thomas: But you are the one who won't let me

214
push the continuum above 5x10 '^^K; indeed most of

the arguments here favor something closer to 3 to

4x10 '^^K. You can't produce a radiation-induced

recombination spectrxim involving N V and O VI from

the radiation field produced by a photosphere at

Tg < 5x10 "^"K. So you are forced to adopt a chromosphere-corona type model in which mechanical heating

is superposed on a continuous radiation field, and

in which Te increases outward through some part of

the atmosphere and reaches values of at least 2 to

5x10^ °K. If we believe the Shapley-Kopal densities

for V444 Cyg and Hanbury Brown's estimate of radii

of the line-producing regions, Te reaches these

values in regions where the electron densities are

10 1^ to 10 ^\

Undevhill : I said the apparent color temperature in the visual was 3 to 4x10** °K. If you had a continuiom at 10^ °K, how would it differ in relative

energy distribution over that short wavelength range

from what we showed here? And if you construct a photosphere at 5x10 with the densities we believe

are correct, you are going to need an enormously

high gravity to hold it together.

Thomas: A photosphere is defined as t=1 in

the continuum; if were 5x10^ °K at such a depth,

I would see it in some way. I don't. And in all

the models to which you constantly refer, you have

electron temperatures of < 5x10 '^^K. This is not

some color temperature but a kinetic temperature,

a parameter you use to compute gas pressure.

So the atmospheric configuration we have to

investigate is clear; maybe we don't know how it got

that way, but we have rough estimates of what it is:

Tg is between 5x10'' °K and 10 and increases outward through at least part of the atmosphere. The

run of density is much less certain, but ne is cer- tainly much higher in the region of line formation

than in the solar chromosphere-corona. So you probably want to consider n^ < 10^** in the region of

line formation and considerably higher values where

the continuum is formed. Whether the He I spectrum

can give information on dilution effects, and hence

on ne, at still greater heights in the atmosphere

remains to be demonstrated. But unless something

intervenes, the opacities in all the lines considered

will be quite high, and solutions of the transfer

equation will be required in order to discuss line

profiles. But of course something does intervene,

as evidenced by point (2) , the breadth of the emission lines: there are large velocity fields. We

cannot be sure what they do to the opacity because

215
so far we haven't been able to decide between a simple expanding atmosphere, a differentially ex- panding atmosphere, and an atmosphere with random

motions. So the opacity is one of the things to

be determined, and it is directly related to the

distribution of velocities over the atmosphere.

The problem we face in investigating the WR

atmosphere is essentially the same as the one we faced some years ago, and still face today, in

studying the outer solar atmosphere: what distribution of Tq and n^ will give something like the

observed spectrum? The added complication for the

WR stars is the velocity distributions - macroscopic

and microscopic. The added simplification is the

increase in our store of knowledge and experience

in how to approach a non-LTE situation. If we could

neglect the effect of velocity fields on the opacity,

the problem would be straightforward: we would assiame some distribution of Te and ne and we would

compute the spectrum. The techniques are well known;

they are trivial for the thin atmosphere and some- what more difficult for the thicker atmosphere. When

we take into account the effect of velocity fields

on the opacity, we require solutions of transfer

problems in moving media; both the opacity and source terms in the source function must be coupled

to the velocity field. It is this kind of approach,

rather than one based on the standard atmospheric

models, that we should be considering here. After

we have carried out these quasi-empirical investigations on the kinds of distributions of Tg , Uq, and

V that might produce the observed spectrxim, we can proceed to investigate the energy balance. This is

the stage we have now reached in the outer solar

atmosphere. In applying the standard atmospheric

models, you have assumed at the outset that you know

the answer to the last question.

Pecker-Wimel : You must make some assumption

about the mechanism of excitation and about T^ in

the continuum.

Thomas: The first is not necessary; all likely

processes can easily be included. I agree with the

necessity to assume something about T^, but I presumed this could be obtained from observations in

the visual and rocket UV.

Undevhill : Bappu gives arguments to justify

his assumption that the non-LTE factors bj. satisfy

the relation b^/bj = 1. And actually the important

parts of Tr are in the region X < 22 8 A where they

cannot be observed.

Thomas: It is unnecessary to assume anything

216
about the bk's; I would not trust any a priori as- sumption. Agreed on Tr, but we seem to be gaining

rapidly in empirical knowledge; let's give NASA a little more time and we'll have the answer. Peoker-Wimel : I am working on such problems

now, trying to include the velocity field. But I think many of your arguments hold only for the 2-

level atom, and that the procedure becomes difficult

for more complicated atoms. It is very difficult to

separate the various excitation processes, especially

when you have mechanical heating and are uncertain

about the continuous radiation field.

Thomas: Well, you and I have a couple of papers

where we show how, in principle, this can be done.

It necessitates the ability to identify the important processes populating any given level of the

atom. A number of people seem to be progressing

rapidly with numerical solutions to these problems

,

mostly without velocity fields, it is true. But I do think we now have an insight we did not have some years ago into the physical effects of the various

parts of the source function: the scattering term

and the source and sink terms.

Underhill : You can't really do this because

you don't know the cross sections with the necessary

accuracy. And furthermore, you have these particular processes of which I have managed to isolate two

or three. But there is no guarantee there aren't

many more of them as yet unrecognized.

Thomas: Our experience is to the contrary, but

we won't know until we try. The problem in which we have the least experience is that involving moving

media. Charlotte Pecker-Wimel has been working on

it; so has Dave Hummer; possibly we could ask them

to outline their thinking.

Pecker-Wimel : Most of my results are still on the machine; I had hoped to bring them along, but

they were not ready. I would only emphasize that

you cannot choose, a priori, a distribution n^ (r)

without being sure that it is consistent with v(r)

Hummer: From the work George Rybicki and I have been doing on line formation in expanding atmospheres, I have gained some insight into how the

different kinds of profiles can be interpreted. Our

work involves numerical solution of the non-LTE line

transfer problem with noncoherent scattering in plane

parallel slabs in which radial velocity, temperature,

density or any other atmospheric parameter can vary

arbitrarily with depth. Since I have had planetary

nebulae primarily in mind in this work, we have always taken the mid-plane of the slab to be stationary

217
with respect to an observer outside. The velocity

on the near side increases toward the observer, while

that on the other side increases away from the observer. We find that the usual double-peaked intensity profile is distorted by a reduction of the intensity on the blue side, so that an apparent redshift occurs. This is, of course, the effect first

obtained by Chandrasekhar and later by many others

for absorption lines; the intensity in the blue is

reduced, leading to an apparently blue-shifted absorption line. The explanation is quite simple: one sees to greater depths in the red than in the

blue, and for the situations considered so far the source function is greater at large depths than near

the surface. An interesting consequence of this ex- planation is that if one had a strong enough temperature increase toward the observer, again with the

velocity increasing in the same direction, one should

see blue-shifted emission and red-shifted absorption

lines. It is worth emphasizing that the red shift

has nothing to do with the receding rear part of the

slab, for even when this part is stationary, the same type of profiles are obtained.

The importance of this work for the theory of

WR stars seems to me to be two-fold. First of all, we can develop some feeling for the kinds of profiles that are associated with various types of

atmospheric motion, including optical depth effects.

Second, when definite atmospheric models are proposed, we can test them to see if the correct line

profiles are indeed observed. I would like to take

spectra of the type shown us by Anne Underhill and, on the basis of the intuition we have now developed,

pick out those lines which are formed by a radiative

mechanism in a turbulent situation and, in the case

of the flat- topped emission lines, those which are formed in a thin atmosphere situation. I think it

is clear that these latter arise because of the large

velocity gradients. The point is that it is now possible to sit down and make a good guess at a kinematical picture for each line.

Smith: I have tracings of the spectrum of

HD50896 (WN5) that show a definite central reversal

in the He II A.4686 line. Figure 8 shows the asymmetry of the profile.

Eumm.er: That is exactly what you would expect.

What would be a central self-reversal in a static

atmosphere becomes distorted in an expanding atmosphere. The blue wing simply drops away.

Underhill : The. great problem in interpreting

these things is that it is difficult to be sure from

218
Figure 8. The line profile of He. II A4686 in

HD50896 from a 9-minute exposure at 16 A/mm taken 26

September 1967, with the Coude spectrograph of the

120-inch telescope of Lick Observatory. All other

strong lines on this spectrogram show asymmetric

profiles similar to that of A4686 but with less

pronounced maxima on the redward side. This is a transient phenomenon; a plate from the previous

night shows normal Gaussian profiles.

your observations whether they are real effects, or whether the tops of the profiles are simply grass.

In photographic photometry you cannot get away from

grass being about 3% of the intensity.

Smith: These are photographic observations,

but the grass is significantly less than the depth

of the self-reversal. Nor is there a possibility of

confusion resulting from a blend with N III X4640.

The center of that line would lie significantly

outside the X4686 profile, and ic is not present.

The exposure time was about 10 min at 16 A/mm. The

spectrum looks like that of a single star, but Kuhi

thinks it may be a binary. I have seen one other

profile like this, a C III line in Campbell's star.

It was slightly asymmetric but not so much so as this. On the whole these profiles are quite rare. Underhill : 1 would emphasize that the main

problem is to be sure you have short exposure, high

resolution spectra. There is a note by Hutchings in

219
Observatory on photoelectric scans of Ha and H3 in

Be stars. Of stars and B supergiants. They see definite changes in profile which you will not see photographically because you are integrating over a longer period of time.

Payne-Gaposahkin : It is worth noting that some novae emission profiles exhibit what looks like

self-reversal until you notice that the shape is

shared by the forbidden lines. This makes it un- likely that the phenomenon is really self-reversal.

Thomas : Charlotte Pecker-Wimel and I had a note in Ap . J. a couple years ago on this phenomenon

in the [Fe X] red line. I agree it is hard to imagine that such effects come from self-absorption,

but we managed to convince ourselves that this was

the explanation. There was a suggestion some years

ago that novae ejected two globs of material symmetrically, but we gave reasons for excluding this

interpretation

Sakade : I think it is dangerous to base any

conclusions on the X4686 line. As we have heard,

it is a rather complicated line in binaries. I think

we must look for other lines in which to study selfreversal effects.

Thomas: That may be true, but when you have

an observation that leads directly to some physics,

you try your best to use it. The existence of a central self-reversal in an emission line tells

you immediately that the optical thickness of the

layer producing the line exceeds about 10. Asymmetry in such a line tells you immediately that

you have a velocity gradient.

Alter: We have known for some 25 years that

most of these lines are strong lines formed in an optically thick atmosphere.

Thomas: I am not aware that this is well known.

What we have known is that the ratios of the line

intensities are not the ratios of the f-values . But

these are two quite separate things, and the one does

not imply the other.

Underhill : We definitely know these atmospheres are thick. When you have such strong lines

in atmospheres of such small dimensions, these atmospheres must be optically thick. O. C. Wilson showed

in 1942 that the dimensions of these atmospheres were small. I suggest He II A4200 is a better line

than A4686 for diagnosing the atmosphere, because it

is not particularly strong, whereas X4686 is probably formed outside the main part of the atmosphere.

Thomas : Everything depends upon the velocity

distribution. If you take 10 R® for the radius of

220
the line-producing atmosphere and a particle con- centration of 10^ , you get a line-of-sight particle

concentration of 10^ . Even if the absorption coefficients were small, this would give a high opacity

if there were no differential velocity. If there is

a differential velocity you must compute the opacities

very carefully. I see no sound basis for your con- clusions, and this is why I lay such stress on things

like the self -reversed cores, which do provide a real, physical basis for inferring something about

the opacity.

West: I would like to describe a microscopic

calculation of the type we have just been discussing, which has been made by Stecher and myself and

is aimed at interpreting the C III spectrum in the

rocket-UV observations of y Vel presented earlier

in this symposium by Stecher. The calculations are preliminary. By making various simplifying assumptions we are trying to get some feeling for what

physical effects are important. This should be re- garded as only a progress report.

The strongest emission line in the spectrum

was identified as the C III intercombination line

X1909. According to Code and Bless (1964, Ap . J.^

139 , 787) , the C III ion has the most completely

developed emission spectriam in this WC8 star. It

would not be surprising, therefore, to find strong

C III lines in the UV. However, the X1909 line is

the transition from the ground state of the triplets

2p to the ground state of the ion, 2s^ ^S°.

Garstang and Shamey (1967, Ap . J.^ 148 , 665) give

the oscillator strength as gf = 3.1x10"^ or A = 190

sec""^. This low transition probability contrasts

with the observed equivalent width of about 14 A

(minimum value, uncorrected for instrumental profile) to provide an interesting problem in the mechanism of line formation.

In order to discuss quantitatively the physical

processes that may be operative, we need a rough

model for the emitting region of the envelope. We assume the WR star to have a photospheric radius of

5 solar radii, and we consider the line to be formed

in a moving envelope whose outer radius is 5 times

the photospheric one. (These assumptions are con- sistent with the results reported at this symposium

by Hanbury Brown.) We assume the velocity field in

the envelope to consist of a spherically symmetric

radial expansion, with zero velocity at the lower

boundary of the envelope (which coincides with the

outer edge of the photosphere) and a linear outward

increase reaching 1500 km/sec at the outer boundary

221
of the envelope. We assume the only random motions

in the moving envelope to be thermal, and we assume

Te = 5x10 '^^K everywhere in the envelope. We assume

the density to decrease outward from the lower boundary as r~^, where r is the radius. We explore a range of values for the density at the base of the

moving envelope. Finally, for the continuum radiation field at the top of the photosphere, we take the

results from Mihalas ' continuum model of a 4x10**

star (1965, Ap . J. Suppl,^ IX, 321). We assume there

is no continuous opacity in the envelope.

Given this simple model, we could now proceed

to investigate the line-formation process by solving

simultaneously the equations of statistical equilibrium and of radiative transfer to obtain values for

the source function and the line profile throughout

the atmosphere. Such a procedure would be more com- plicated, for two reasons, than those which now exist

in the literature. First, even in this simple model,

there are very strong differential velocity fields.

Second, the atmosphere departs severely from a planeparallel model. As mentioned, this present investigation is preliminary and exploratory, so we proceed to make three further simplifications.

First, we assume the total line profile to be

the superposition of profiles from a number of atmospheric elements, each of half-width equal to one thermal Doppler half-width, and each centered at a frequency corresponding to the line-of-sight velocity

with respect to the observer. Because of the velocity gradient, the contributions from other regions

of the atmosphere will be shifted out of the frequency band emitted by the atmospheric element con- sidered. The dimensions of the contributing atmospheric region are fixed by:

Z2 - zi = 2u/(dv/dz) = constant (1)

because of the assumptions of an isothermal envelope

and a linear velocity gradient. The line-of-sight

coordinate is z, u is the thermal velocity, and v is

the mass-motion velocity.

Second, when we observe an emission line, we must decide whether it arises wholly from a geometrical effect because the envelope is larger than the

photosphere, or whether it is intrinsic and arises

because the envelope between the observer and photosphere emits more in the line than it absorbs from

the continuum at the same wavelength. In the present simplified calculation, we consider the question of intrinsic emission. Thus we restrict our

222
attention to the emission from that part of the en- velope lying between us and the photospheric disk.

We will then underestimate the energy in the emission line, relative to the observations which refer

to the entire envelope.

Third, rather than computing the source function from the equation of radiative transfer, we proceed by assigning it various values and computing

the effect on the emergent radiation. We assume a v-independent source function, so that we obtain its

value by assuming the population ratio of upper and

lower levels.

Under these conditions the equation of radiative transfer along the line-of-sight from the stellar disk through the envelope to the observer becomes : dl 2 1 dT2 1 where

l2i - S21 (2)

S21 = E21/K12

(3)

E21 = N2A2ich/8TTU (4)

K12 = (N1B12 - N2B2 1) ch/8TTU (5)

where u is the thermal velocity of carbon; Nj is the

occupation number of the level j; and the A21, B12?

and B21 are the Milne form of the Einstein transition probabilities.

Equation (2) integrates to give:

I21 = I21 (t) e +

T S21 e dx , (6)

0 and under the assumption of a constant value for the source function, S21, equation (6) becomes:

I21 = l2i(T) e-^ + S2i(l - e-'') . (7)

We assume there is no spectral line at the top of

the photosphere, so I^i is just the continuous ra- diation field given by Mihalas ' computations. Thus

the quantity of interest is I21/I21/ the ratio of

223
the emission in the line to that in the continuum.

Clearly, from the restrictions set by the assumptions , we will get only the center and the blue

part of the line profile; the red part will be oc- culted by the photospheric disk. In Figure 9 we plot the value of I21/I21 ^'^ line-center for

various assumed values of (N1+N2) at the base of the

envelope and for various assumed values of N2/N1.

The value of x follows directly from equation (1)

and these two assiamptions on the values of the Nj :

T = K21 2u/(dv/dz) . (8)

For convenience, we also give the values for b2/bi

,

the ratio of the departures from a Boltzmann distribution.

The observed quantity is the ratio of total

flux in the line to that in the continuum, Fgi/Ffi;

so as mentioned earlier the observed and computed

quantities are not strictly comparable. Let us,

however, ask what conclusion we could draw if they were comparable . The observed value for F21/F21 is at least 4 and could increase to 7 for an infinitely narrow

line corrected for instrumental profile. Taking the

conventional H/C abundance ratio and assuming all C occurs as C^"*" , the curve in Figure 9 for (N1+N2) = 3.2xi0^/cm'' represents an H density of 10^ ^/cm^ in

the envelope. So if this curve indeed represents

the conditions in the envelope, we require N2/N1 >. 14 to reproduce the minimum observed ratio of line

to continuum. Such a value for N2/N1 requires a population inversion of the two energy levels; the

absorption coefficient becomes negative; and a laser action will amplify the radiation in the line.

If such a situation exists, it is indeed very interesting. Its existence depends upon the large N2/N1

ratio; the next problem would be to explain it.

To explain the observations other than by a large N2/N1 ratio requires that one increase by a factor of 5 either the H density at the base of the

envelope or the C abundance ratio - or that one change the underlying assumptions. While the C

abundance remains open to controversy in WR stars

,

10 ^^/cm^ is thought to be an upper limit to the

electron density because collisional de-excitation

would dominate spontaneous emission at higher densities . Of the assumptions underlying the theory, that

which allows I21/I21 to be compared with F21/F21 is

224
_l I I I I I I I I I I I

I 2 3 4 5 6 7 8 9 10 II 12 bz/bi

Figure 9. The ratio of the emission in the

line to that in the continuum at the line center for

various assumed values of (Ni + N2) at the base of

the envelope.

the most open to question. We are now in the process

of refining our calculations to eliminate the necessity for this comparison. Our present point is

simply that these results have raised the interesting possibility of laser action, and that they should

be followed up by more detailed calculations.

For reference, I should mention that we used

the following values in our computations:

I°i = 4.78x10"^ ergs cm"^ sec"^ v"^

225
u = 1.0x10^ cm sec~^

2u/(dv/dz) = 1.9x10^° cm . Hummer: What is it that makes this transition

alone show laser action? How does it differ from

the other similar transitions?

West: It is reasonable to expect that you

could get an overpopulation in a metastable level

because the ways of getting out of it are few and

slow.

Hummer: But there are a lot of other atoms

with similar metastable states.

Underhill: Not atoms with abundant ions that

give you observable lines.

Steoher: On the observational side, X1175,

which arises from the level, is very strong.

It is seen in emission and absorption, Doppler

shifted about 5 A. This indicates a considerable

population of the triplets in the envelope. The

collisional depopulation comes at about 10^^ electrons/cm^. We don't have an exact value since we

do not know the cross section.

Underhill: You do have to maintain an appreciable population in that metastable level of C III.

It is the upper level for the A1909 transition, but

it is also the lower level for the A4650 absorption

line, which is one of the most prominent features in

the spectrum. This makes me think the densities

should be low; I also believe 10^^ is an absolute

upper limit.

Steoher: I might add that there are other data

for a number of other recombination lines. We have

not yet analyzed it, but I think a number of interesting things may develop when we do. In particular

we find N IV at X1488 and 0 V at A1216. There is

the possibility that the O V line is exciting He II

A4686.

Underhill: I was hoping you had data on the

O V lines at A2750.

Steoher: There is only one WR star we can ob- serve with rockets; we need a satellite to observe

the others.

Aller : These WR stars seem to provide such a formidable array of difficult problems, all of which

must be solved simultaneously, that I wonder if it

would not be wise to select some simpler objects

that exhibit at least some of the spectral features

observed in WR objects and see whether we might

226
learn something about the physical processes involved. I have in mind the P Cygni stars, which

are similar to the WR stars, but with simpler line

profiles; the profiles have rounded emission com- ponents with violet absorption edges. It is nor- mally thought that the P Cygni stars and the WR

stars can be readily differentiated from one another.

Beals has studied this problem; I wonder if he

would like to comment.

Beals: I was particularly interested by Miss

Underhill's mention of a number of stars with nar- row lines. These are the WR stars of type WN7,

and I believe that HD151932 is one of the group.

But had we not defined the category of WR stars, I feel sure these objects would have been classified

as P Cygni. Indeed I am not sure that they should

not have been classified as P Cygni. I am inclined

to say that there is no sharp line between WR and P

Cygni stars. I agree that the P Cygni phenomenon

may shed light on the WR phenomenon, and I feel

that the two should certainly be studied together.

There are 50 or more of them altogether.

Alter: Two points should be mentioned; First,

the excitation level in the P Cygni stars is gen-^^

erally lower than in the WR stars , so the theoretical treatment may be easier. Second, some objects

are observed to flop back and forth between a P

Cygni spectrum and something else; I refer in particular to S Doradus, which every now and then has

a nice P Cygni spectrum. Such variation makes things

less agreeable.

Underhill : Mart de Groot, at Utrecht, is working on P Cygni profiles for his thesis. This work

won't be available for about a year, so I might just

summarize the highlights. P Cygni itself is classified as Bl q and is of a much lower level of excitation than the WR stars; the hydrogen lines for ex- ample stand out in its spectrum. In those line

profiles associated with expanding shells, the emission peaks are less striking than in a WR spectrum,

and the absorption peaks dominate the line. I think

that whether the line is dominated by emission or absorption peaks depends very much upon the level

of excitation in the shell. In the hydrogen lines

you get a sharp emission peak and a strong absorption peak, sometimes two absorption peaks. The

amount of the absorption varies with time, while I believe the emission peak stays fairly constant.

In the He I lines you get much more absorption and

hardly any emission. You do find a few lines with

only emission, which just means that there are not

227
enough atoms in the lower levels to give absorption

in the shells; most of the lines you observe are subordinate lines.

In P Cygni spectra, there are no double emission peaks such as you get in Be stars. P Cygni

stars give the feeling of a spherical expanding atmosphere; like WR stars they give the impression of

spherical symmetry. Be stars and novae, on the other hand, give the feeling that things happen in a plane, that they are more nearly cylindrically symmetric.

Hummer: What characteristics of the spectra

tell you that?

Undevhill : The only way of getting two separate

emission pips is from a rotating star in which material is concentrated in a plane.

Hummer: But this depends on the amount of material present; I can see your point if the material

is only marginally thick.

Underhill : Be stars and novae shells are only

marginally thick; shell stars are thicker.

Beats: The complex of lines, which apparently

exists in P Cygni as a temporary phenomenon, seems

to be permanent in HD190073. It is one of the most

fascinating lines I've seen in any star, and I think

it is changing in detail but not in general outline.

It is an undisplaced emission line with a central

absorption minimum and two strong absorption lines

on the violet edge. You frequently get this sort of

thing in novae, where it is easily explained in terms

of two moving shells. In this star, however, it has

been like this for at least 30 years and probably

much longer. HD19 0073 is between 7th and 8th magnitude; I'm delighted to have a chance to mention it

here in the hope that someone with modern equipment

will observe it again.

Pecker-Wimel : Is this profile observed in many

lines or in only one?

Beats: There are profiles available for the

hydrogen lines, and they show three absorption com- ponents corresponding in velocity to those of Ca II.

The hydrogen lines don't look at all like the Ca II

lines, and although the velocities are roughly the

same, the various components of the line do not have

the same relative intensities.

Johnson: From what you said about the very

great life of this double absorption feature, I gather you don't necessarily believe it can be ex- plained wholly by a classical expanding shell. Are

you suggesting that the doubleness of the feature

implies something unique?

228
Beats: I tried to explain the profile in terms

of the acceleration of atoms. I assumed that the

large-velocity component arises from atoms close to

the star and that as the material moves out, it ex- periences different rates of deceleration. When it

is absorbing strongly, it is moving outwards with a uniform velocity or with a very small deceleration;

when the absorption is weaker, the material is decelerating more rapidly; and when it reaches zero velocity, there is a shell surrounding the star at some great distance, showing an emission line with

a central absorption. This was the only model I could think of; I am not particularly happy with it;

if I really believed it, I would not be making an extra effort to get other people to study the star

and explain the profile.

Alter: FU Orionis is a star that suddenly

brightened and evolved a P Cygni spectrum that

has just remained ever since.

Undevhitt: Deals' discussion points out that

the major difference between the novae on the one hand and the WR and P Cygni stars on the other is

that the latter show a spectriam constant in time.

The displaced absorption component indicates you

are seeing an outward flow of gas, but you are not

always seeing the same atoms; you are seeing the

gas at a particular geometrical point where the

density and excitation conditions are such as to

produce the line. You have a steady-state velocity

flow. Novae, on the other hand, do not reach equilibrium; they may go through a phase showing the same spectrum as WR stars, but it is a transient

phase. A similar, possibly intermediate situation

holds for the shell stars. A shell star may give

evidence of a steady-state outward streaming for a long time and then suddenly change; you will then

have evidence of both a stationary and an outward

moving shell. In 48 Librae, for example, the shell

lines have given a radial velocity variation with a period of apparently 10 years. But recently the

star has gone crazy and shown negative velocities.

I took a spectrum at Kitt Peak that showed double

Na D lines: a strong, rather sharp line and a fuzzy patch to the violet; you obviously have two

pieces of expanding gas.

Attev: Doesn't the supergiant A star, a Cyg

,

show a P Cygni profile for Ha? The line spectrum

is otherwise constant in time , although the atmosphere itself shows oscillations.

Beats: I think the P Cygni characteristics of

Ha in a Cyg are somewhat variable.

229
Nariai: Kikuchi has taken spectra of Ha in 3 Orionis and finds that its profile varies with a period of about 20 days.

Sahild: This raises a question I hoped to ask

of Deals. From looking at a fair number of spectra

of very luminous stars, it is my impression that

many of the 50 or so stars he identified as P Cyg

type have spectra that are fairly normal for highluminosity stars except that they show a P Cyg profile at Ha and perhaps a few of the higher Balmer

lines. Could you tell me how many other spectra

you know that show P Cygni profiles in He and other

lines besides H?

Beats: I would say that the number of stars

showing spectra similar to P Cyg right from one end

of the spectrum to the other is so small as to be

negligible. I think your remark about P Cyg characteristics being imposed on normal spectra is probably

correct. But although we associate P Cyg type spectra

with rather high luminosity stars, I don't think all

the P Cygni stars are highly luminous ; indeed I think

some are not highly luminous at all.

Sahild: Just to be very clear, let me amplify

my last remark. There are one or two highly lumi- nous stars which do have really abnormal spectra;

P Cygni is not unique. You have already mentioned

the two stars, HD151804 and 152408, in the association SCO OBI. These do have very pronounced spectral peculiarities as well as peculiar P Cygni profiles in the hydrogen lines.

Payne-Gaposchkin : The companion of o Ceti is

a very low luminosity star, visible only when the

long period variable is faint, yet it is a kind of

P Cygni star - a bright line, high temperature star.

Also, is there not a suggestion somewhere in the

literature that P Cygni is an eclipsing binary?

Underhilt : Two Russian astronomers reported

in the Variable Star Bulletin that they had found

a half-day period in the luminosity of P Cygni and

suggested it might be a W Ursa Majoris star. de

Groot has measured radial velocities on all his

spectra, and can show that they do not follow a halfday period. He found velocity variations of much

longer period in some of the hydrogen lines; we think

this is an atmospheric oscillation.

Peoker-Wimel : We have from time to time spoken

of the similarities between Of and WN stars. Is

there any evidence of fast rotation in WN stars?

Underhill : Rotation can be determined only

from observations of absorption lines that exhibit

the well-known dish shapes. We have not detected

230
any such absorption lines. Therefore I know of no way of detecting rotation.

Pecker-Wimel : I have heard rather vaguely of one attempt. When WN stars and Of stars in the Large

Magellanic Cloud were plotted on the HR diagram, the

WN stars fell in that region of the diagram where

rotation was observed. So I am wondering whether

you might be getting a blend of rotation and expansion in the observed WR profiles. The profiles com- puted by Sobelev included both rotation and expansion

and showed that you could get a central dip in such

an emission line.

Underhill : Yes, if you have an optically thin

atmosphere that is extended. A typical Be star has

a shell, and the atmosphere is spinning, so you get

two separate emission lines. But I don't understand

the investigation you reported; I don't believe we

know that the Of stars are rotating. It is true that

some of them have very broad lines, and when we have

an absorption line 4 A wide with a 10 percent central absorption, it is a nice point to prove that

the broadening is due to rotation and not to microor macroturbulence . All the observed lines from 0-

stars have this characteristic shape; they indicate

rotational velocities of less than 200 to 300 km/sec.

But again the interpretation depends very much on what you assume for the source function in the atmosphere. The usual procedure, which we know is poor,

is to assume Kirchoff's law and to match the line

profile by introducing microturbulence as a fudge

factor. To be more realistic, we should introduce

scattering into the source function. This would

give us a broad, strong line for considerably smaller

rotational velocities. So until we can separate

the detailed physical processes producing the line,

I think it is hard to separate the various kinds of

velocity fields.

Limber: Now that rotation has been introduced,

I should like to comment on my model of a few years

ago in which I tried to interpret the WR phenoitjenon in

terms of a rotational forced ejection. Though I am not completely convinced of the model, I am not yet

ready to abandon it. I do not intend to discuss the

suggestion in detail; you can read th^ literature

(1964, Ap, J., 1^9, 1251). I would like however to

draw your attention to several characteristics of

the WR phenomenon which, it seems to me, are crucial

to our understanding any model of a WR star. Several of these characteristics involve questions of

rotation and rotational effects, and they may clarify the importance of such effects both in the model

231
of the star at the WR stage and in the evolutionary

history leading to that stage.

First there is the simple kinetic picture of

the WR star. Our first-generation model by Beals

consisted of a very extended, expanding atmosphere.

This simple picture had to be changed for two reasons: First, the absence of a transit-time effect

(phase difference in time of eclipse as observed

photometrically and as observed in the velocity

curves) meant that the region of line emission could

not exceed the region of continuum emission by more

than a factor of about 5 in radius. Second, the

occultation effect (asymmetric line profile) , which

should be present in a model of simple expansion,

was also absent. So we have Kuhi ' s second-generation model, which consists of an inner chaotic re- gion with random velocities of several thousand

km/sec, where the strong emission lines are produced,

surmounted by the low-density expanding region. Now

I would point out an alternative second-generation

model that is equally consistent with the absence

of transit-time and occultation effects. Here the

chaotic velocities in the inner region are replaced

by rotational velocities that decrease outward. So we must make a decision between these two alternative

kinetic models, both of which provide the differential motions necessary to explain the line widths.

There are two kinds of evidence. One is the

geometry of the emitting regions; the other is the

binary character of many of these stars. The geometry, which involves the distribution of mass as well as of excitation, is determined by the rotational instability which in turn is affected by the

presence of a companion. Consider the single star.

At speeds such that the centrifugal and gravitational forces on a mass element at the equator are equal, rotation can readily move mass into a circumstellar envelope, given an appropriate viscous

agent. Thus the centrifugal acceleration provides

support for the extended atmosphere. Moreover, if

the rotation is differential, as may be the case in

the photosphere and must be the case in the envelope, the velocity gradients can provide significant

mechanical energy for exciting the envelope. Thus,

for example, it is possible that some of the chromospheric effects we find in main-sequence F- and Gstars result from differential rotation. The cor- relation that O. C. Wilson finds between chromospheric activity and stellar age may well result

from the reduction with time of differential rotation and thus of energy for chromospheric activity.

232
Rotation seems generally to be a source of instability in stellar evolution. A star contracting to

the main sequence may well face a rotational crisis.

In early post-main sequence evolution it appears

that the Crampin-Hoyle mechanism can bring rapidly

rotating stars to another crisis, thus producing

Be stars. In later post-main sequence evolution,

as Paczynski and others have suggested, contraction

may again bring at least some classes of stars to

a rotational-instability crisis.

Now consider the binary aspect. Here we have

another source of instability which produces mass

loss, as has been discussed. I think one very important piece of evidence is the rather sharp

boundary to the region where emission takes place.

In eclipsing systems, the WR component appears to

be about twice as large as an eclipsing object as an eclipsed. In the system V444 Cyg , from which

the Krons drew this conclusion, I note that the

boundary of this inner emitting envelope lies very

close to that of the inner Lagrangian surface about

the WR component. This may be coincidence, and the

boundary of the emitting region may not be as sharp

as the Krons suggest, but it is an interesting fact.

It suggests that the inner Lagrangian boundary may

play a significant role in fixing the boundary of

the emitting region. This can be true only if the

velocities in the vicinity of the inner Lagrangian

boundary are small relative to a frame rotating with

the orbital motion. So there should be no very large

velocities in the vicinity of the inner Lagrangian

boundary

Thomas : So you are suggesting that if in the

region of line emission there are any high-velocity

chaotic motions or any large differential rotation

velocities, they must vanish at the outer edge of

this region. Then a large expansion velocity cannot

have been reached at this boundary.

Limber: Right. This implies that the velocity

of expansion is not large here, nor are there any

random motions with large velocities.

It has been suggested that WR stars in close

binaries have narrower lines than single WR stars.

One might speculate that less kinetic activity is

required in order to get mass ejection from a WR

star in a close binary than from a single WR star.

It would be easier to ease matter through the inner

Lagrangian point for the binary than to expel it

from a single star.

Thomas: You mean that when a WR star has a companion, it somehow knows it doesn't have to have

233
as much internal instability, or whatever produces

the "chaotic shell", as when it is single?

Limber: What the mechanism of instability is,

I do not know. But I suggest that if the WR star

is a member of a close binary, it can produce the

required mass flux with less instability. I recall

Miss Smith mentioning that WR stars in close binaries

are not associated with ring nebulae. One might

infer from this that any mass lost by a WR star in a close binary is transferred to the companion and not

lost to the system.

Kandel: Thomas has asked how the WR star knows

it doesn't have to be as unstable when it has a com- panion. Possibly it gets this information through

the boundary conditions . Suppose you have some sort

of overstability as the exciting mechanism; then the

amplitude at which everything becomes steady may be

considerably lower in binary systems when you have

a companion to siphon off the mass. In practice the

boundary condition might impose a certain flux of mass

and momentum at a given radius from the star. This

radius would be much smaller for a binary system.

Limber: Let me comment on the observed longitude dependence of the spectroscopic properties of

WR stars in binary systems. Hiltner found twenty

years ago that in CQ Ceph, He II A4686 is strongest

at primary and secondary eclipse and weakest at the

two elongations. This suggests several alternatives

for the distribution of the emitting material; in all

of them it is spherically asymmetric and located in

a region where it is influenced by both stars. Such

effects suggest that the geometries of the excitation and velocity fields are fixed with respect to

a frame rotating with the orbital motion. The effects of the 0-type companion are not necessarily

overpowered by those of the WR star. If this logic

is substantiated, we may be able to rule out some possible excitation mechanisms and geometries.

Finally, I would like to remark on a point that

does not deal directly with rotation or binary character: the relation between other objects and WR stars,

with respect to our models of WR stars. Consider the

Of stars and the nuclei of planetary nebulae. None

of us believes the Of stars represent a late stage

of evolution accompanied by pulsational instability

in a star that is essentially pure helium or carbon.

If the relation between the Of and the WR stars is

superficial, we have no problem; but if the Of

phenomenon is essentially the same as the WR phenomenon, we should worry. The same holds for the nuclei

of planetary nebulae. If the binary characteristic

234
is vital to the WR phenomenon, and if planetary nuc-- lei are not binary, it is awkward to identify one with the other. These difficulties cause me to believe that in the models so far proposed, we do not

understand the WR phenomenon. I believe the possibility of rotation should not be ignored. At the same time, I am certain that the details are con- siderably more complicated than I suggested several

years ago.

Underhill : I would feel much better if you had

been more explicit about those possibilities for

excitation and geometry that you say may be

ruled out by the effects you summarized. And

I would also caution against pushing too far such

comments as yours on the Of stars. I believe the

apparent relation of Of stars to WR stars is simply

a spectroscopic chance. In Of stars we observe many

of the lines in absorption that we observe only in

emission in the WR stars. We can guess the excitation temperatures required to form these lines in

absorption, and this is in fact part of our basis

for saying the electron temperatures in WR stars are

3 to 4x10** °K. But before using the properties of one spectrum to reject a model for the other, we must be aware of the differences as well as the

similarities between the two.

Nariai: If we are discussing general models

for mass ejection, I would like to suggest radiation pressure as a possible mechanism. The equation

of motion for one-dimensional stationary flow, including a radiation field, is:

v(dv/dr) = ^K^F^/c)dv -p-MdPg/dr) - GM/r=

where v is velocity; and are the monochromatic

absorption coefficient and the radiation flux re- spectively; Pg is the gas pressure; and M is the mass

of the star. If the radiation term dominates, we obtain an order-of-magnitude estimate by setting all

the other terms on the right equal to zero and integrating. We get v 100 km/sec for Tj- = 3x10 "^^^K and R = 10^ km. This is a little smaller than the

observed values of v, but possibly it can be increased if the integration over v is properly treated.

In this mechanism, a part of the radiation flux

is consumed in driving the material up to 1000 km/sec

and in pushing it out of the star's potential. Typical values are L^-^d = 10^®, Lk.e. = 3x10^^, and

Lq = 10^^ erg/sec for M = 5 M©,*r'= 7 R© , Tg = 3x10**^,

235
V = 1000 km/sec and dM/dt = 10"^ M®/year . Lk.e. = 3x10^^ erg/sec corresponds to F^.e. = 10^^ erg/cm^/

sec . We cannot get such a large value other than

by the conversion of radiation into kinetic energy

flux. For example, a maximum acoustical energy

generation from the convective zone is 10® erg/cm^/sec

for normal composition at 6x10 ^°K, and 10^ for a helium rich atmosphere at 1.3xl0'*°K. There is no convection zone due to ionization of He II around

Tg 3x10 '*°K and log g < 4. As pointed out by Miss

Smith, shock waves due to nuclear instability are a fascinating mechanism. They may be able to carry

the required energy, but the flux is reduced by 1/e

in a distance of the order of 10^ to 10^° cm, which

is pretty small compared with the size of the envelope, 10^ cm. Then the velocity is expected to

rise steeply near the surface and to decrease gradually toward outer space, which makes it a little

difficult' to reproduce the observed profile.

236
PART D

A SUMMARY OF PROBLEMS, IDEAS,

AND CONCLUSIONS ON THE PHYSICAL

STRUCTURE OF THE WOLF-RAYET STARS

INTRODUCTORY SPEAKER: Richard N. Thomas

Joint Institute for Laboratory As trophy sios ^ University of Colorado ^ Boulder

I. Introduction 2Z9

II. Summary 241

1. Taxonomy of Classical WR Stars 241

2. Specification of Quasi-WR Objects 245

3. Spectral Characteristics of the WR

Atmosphere 24 5 4. The Structure of WR Objects, and its

Effect on Atmospheric Features 24-8 5. The Structure of Quasi-WR Objects and

its Effect on the Atmospheric

Features 250

DISCUSSION 251

237

I. INTRODUCTION

In each of the three preceding days, we have

summarized and discussed one of the three broad as- pects into which we divided the study of objects

exhibiting WR spectra: taxonomy, spectral features,

and the interpretation of spectral features. Today

we shall try to synthesize these discussions to see what problems we can pose, what ideas have come up,

and what conclusions we can reach on the physical

structure of the Wolf-Rayet stars. I propose to

start this synthesis by outlining what I think has

come out of these discussions. At the beginning of

the symposixam, I presented several questions which

it seemed to me we must answer before we can claim

any real understanding of the physics of these stars.

Naturally my outline here will be guided by these

questions. But by presenting them as an introduction rather than a conclusion, I am giving you the

opportunity to criticize and correct those aspects

which arise from my personal myopia. The object of

such an outline is the same as that of the summaryintroductions on the preceding days - to set up a skeleton containing the form of the subject, which

you can alter and flesh out. Possibly from such

individual views of the last three days, we can produce some kind of multi-stereoscopic picture that

has coherence.

At the beginning of the symposium, I proposed

that within the broad category of objects that exhibit

some variety of WR spectrum, we set up two subcategories: (A) "classical WR stars" whose spectra - the "pure" WR spectra - are sufficiently well defined to be unambiguously identified, and (B) "quasiWR objects" whose spectra resemble the "pure" WR

spectra sufficiently closely to suggest some common characteristics in their atmospheres. I suggested

two alternatives for a physical interpretation of

the classical WR stars. Either they are a distinct

kind of star, whose physical properties lie within

narrow limits, or they represent a particular atmospheric configuration which can occur in many types

of stars and objects with widely differing physical

properties. The first alternative would imply that

the classical WR stars are distinct "WR objects",

with a distinct internal as well as atmospheric

structure. The second would simply describe an atmospheric configuration, independent of internal struc239
ture. The question of "objects" versus "phenomena"

has been discussed extensively in this symposium.

So our first question is:

1. Can we, on' the basis of spectra alone, es- tablish a distinct class whose members will have some or all of their other physical properties lying within some limited range?

Our second question follows immediately from the

first and already implies some measure of physical

insight:

2. Can we identify a class of quasi-WR objects

by certain spectral features , common to them and to

classical WR stars, that imply some common atmospheric conditions?

The first question can be investigated wholly

empirically; to answer the second question, we must

pose a third.

3. What are the unique properties of the classical WR atmosphere that produce the pure WR spectrum?

To answer question (3) we would normally begin

by isolating the class of WR objects or phenomena

and then proceed to investigate the state of their

atmospheres and the mechanisms which produce and

maintain these states. If the classical WR stars

are indeed objects, we would ask: What is unique

about their internal structure that produces such

an atmosphere? If, on the other hand, they are phenomena, we must ask: What do the atmospheres of

such a wide variety of objects have in common that

produces the WR spectrum? But as has been suggested

several times during this symposium, there is another

approach to question (3) . We could study the quasiWR objects and, by comparing them with normal stars

and with the classical WR stars, isolate those properties responsible for the WR features. We could

then try to determine whether these same features were fundamental in producing the pure WR spectrum.

Finally, of course, we would hope that given

the answers to questions (1) to (3) , we would be

able to answer the final two questions in the se- quence of our understanding of the WR stars:

4. What is the complete structure of the WR

object - if indeed such an object exists - and how

does this structure give rise to the atmospheric

features?

5. What are the complete structures of the

quasi-WR objects, and again how do they lead to the

atmospheric features producing the quasi-WR spectrum?

The symposium was implicitly organized on the

logical pattern of these questions. Looking back

240
over the three days, I think we might have been more

explicit about it. But if we now follow this pattern in the summarizing session we may be able to

distinguish some success as well as some glaring

omissions

II. SUMMARY

1, Taxonomy of olassioal WR Stars

During the course of this symposium we have

from time to time heard the plea that we should not

be too dogmatic in our assignment of stars to, or exclusion of stars from, the WR class. This has

been especially true when we have tried to distinguish between pure and quasi-WR spectra. We have

been told that high dispersion spectra show such

individuality that each star effectively defines its own subclass. The implication is that the WR category is not distinct and isolated, but rather that

it lies toward one end of a possibly multidimensional chain, of which its subclasses form a number of

links, but to which sequences of other objects showing quasi-WR spectra also belong.

I would be the last to argue against such a picture, particularly in view of my long-standing

assertion that a WR star simply represents an ex- treme example of an atmosphere fed by mechanical

energy and momentum. But we must not allow a preoccupation with minor differences to divert us from

the basic problems. The general WR class is characterized by features unique with respect to normal

stars and with respect to classical models of stellar atmospheres. The characteristics which define

the subclasses are perhaps less unusual, except for

those spectroscopic features which differentiate the

WC and WN sequences. So it is on the unique spectral features of the general WR class that we must

concentrate in trying to isolate unambiguously those

physical effects which produce them. A first step

in this direction is to inquire into the non-spectral

characteristics of the WR and quasi-WR stars and

their relation to the spectral features.

With respect to WR and quasi-WR objects, we are,

in a sense, in a period of spectroscopic diagnostics

analogous to the era prior to the identification of

the physical parameters that determine the "normal"

spectral sequence. Behind much of our difficulty in

interpreting WR spectra is the implicit assumption

that we know what these parameters are. Over and

241
over again during this symposiiim we have heard such

remarks as: "One of our goals is to arrange WR stars

in a sequence of Teff because T^ff is a parameter

basic to stellar structure and evolution". Again

during comparisons between stellar types and between

various kinds of luminosities we have continuously

evoked a location on the HR diagram. I quote an interchange

Underhilt: ...because the interpretation of

intrinsic colors in terms of effective temperatures

is done by means of model atmospheres. I'm inferring

that because Schild put his stars on an HR diagram

as one would B stars, he was prepared to assign them

effective temperatures. This is the normal procedure when interpreting an HR diagram.

Schild: That is indeed what I had in mind.

May I remind you of the two advances - one theoretical, one empirical - on which the classical

interpretation of stellar spectra is based and of

the order that their adoption brought to the array

of facts on stellar classification. The theoretical advance was Fowler's assertion that the spectral sequence could be understood in terms of thermodynamic equilibrium populations for the excitation and ionization levels, i.e., in terms of temperature and density. The empirical advance was

the relation between luminosity and spectral class,

which amounts to a specification of temperature in

terms of the energy in the radiation field. Thus

the temperature and occupation numbers of all energy states follow directly from the total luminosity. We can construct model atmospheres for each

spectral type, and assuming that luminosity, radius, and gravity suffice to characterize such atmospheres , we can construct the corresponding interi- or models. Conversely, we can try to infer these

parameters from the spectrum and from them infer

the interior models.

By returning to this pre-model era of spectroscopic diagnostics, we are seeking, in our study of

the WR and quasi-WR objects, just that kind of key

provided by the Fowler and the HR relations.

So in this summary of the taxonomy of the WR

objects, I repeat my list of those spectral features

which unambiguously characterize the WR category.

In each debate over the inclusion of a given object,

we find that one or more of these four criteria have

been violated. Each time we try to apply a "standard"

model atmosphere to a WR object, we find one or more

of these features are incompatible with the model.

242
1. The spectriim consists almost wholly of emission lines. When absorption lines occur, they occur

as satellites at the violet edges of the emission

lines

2. The emission lines are very broad. Interpreted as Doppler-broadening , the widths correspond

to differential motions of some hundreds to thousands

of km/sec and are not necessarily the same for all

ions

3. The lines in any one star represent a wide

range of excitation and ionization. The excitation

level of the line spectrum is generally much higher

than that of the continuum as estimated from its

spectral energy distribution.

4. The spectrum falls into one of two groups.

Either it shows strong lines from carbon and oxygen,

or it shows strong lines from nitrogen. Both groups

show strong lines of helium plus other weaker lines.

Now it seems to me that given these four criteria, I already have a basis on which to answer question (3) and to make some inferences on question (4) . Also, I now know what to look for in trying to answer question (2) and to establish a category of quasi Wolf-Rayet objects. But we will come back to these points later.

In our original plan for the symposium, Lindsey

Smith's review paper was to have been aimed at question (1) . She posed the questions somewhat differently. Instead of considering both classical WR

stars and quasi-WR objects, she restricted herself

to the former and considered whether a WR object

exists. She attempted to establish the existence of common physical properties within subclasses and to

clarify the meaning of the various subclasses by

identifying each subclass with an evolutionary stage

of WR objects. She attempted to decide whether these

subclasses represented successive evolutionary stages

of initially similar objects or alternative configurations of objects with somewhat different initial

conditions. In the first alternative, the WR class

as a whole comprises objects of like properties; in

the second, this is true only for the subclasses.

Thus a third possibility, that the subclasses represent different aspects of a general WR phenomenon in

different kinds of objects, was a priori excluded.

In supporting her conclusions, she impinged heavily

on question (4) , the structure of WR objects and its

influence on those atmospheric phenomena that give

rise to the WR spectrum. But I reserve comment on this last aspect of her summary until I have considered questions (2) and (3)

243
We come now to the significance and the homogeneity of membership within the subclasses. It

would be pointless for me to try to summarize the

criteria proposed for the various classification

schemes. Lindsey's differs from the others chiefly

in the absolute luminosities attached to the subclasses, and in the emphasis on the presence or absence of various lines, e.g., 0 VI. You may wish

to push these points further. After listening to

Kuhi ' s remarks, it seems to me that the outstanding

inadequacy lies in our knowledge of the continuum

characteristics as a function of subclass. Both

Lindsey and Kuhi agree on two points: first, that

although there are few data for comparison of the

two sequences , the continuum temperature appears to

be lower in the WC than in the WN class; and second,

that the energy radiated in the lines is greater in

the WC, relative to the continuum, than in the WN

class. Now suppose we were to argue that the mechanism producing the continuous spectrum differs from

that producing the line spectrum. Then these 2 points

could be interpreted as a difference not so much in

the "intensity" of the line-producing mechanism as in the "intensity" of the continuum-producing mechanism. For those stars investigated in detail, the

energy distribution in the continuum seems to depart

significantly from that predicted by the theoretical

main-sequence models and also from that of mainsequence stars. Indeed it is the supergiants which

WR stars seem to resemble most closely. So again we have evidence that the general WR class is internally

homogeneous and that it differs from the class of

normal stars: it does, however, seem to have a closer

relation to certain other exceptional types.

As for the conventional physical properties of mass and luminosity, we again have little evidence;

but what we do have suggests a first-order difference in the masses in the WC and WN sequences, with

the WN the more massive. There is disagreement on the reality of the very brightest magnitudes {"^ -6)

for the WR stars, but no strong evidence for a difference between the two sequences.

In both the WC and WN sequences the luminosities

seem to increase toward later subclasses, whereas the

excitation of the line spectrum seems to increase

toward the earlier classes. So if one assumes that

line excitation correlates with energy in the con- tinuum - and this is not at all obvious - the size

of the object must, as Lindsey suggests, increase

toward the later subclasses. Finally, Lindsey finds

evidence for an association between ring nebulae and

244
certain of the WN subclasses, but no such association with WC objects.

From these data, I end up with the feeling that

there is homogeneity of physical properties within

each of the subclasses, but that these properties

change significantly from one subclass to another.

This conclusion is reinforced by Lindsey's suggestion that the WC subclasses arise from differences

in initial chemical composition while the WN subclasses arise from differences in initial mass. These are points for you to discuss. From the strong

differences of opinion already expressed at various

tim.es during the symposium, I am sure that at least

some of you will have something to say on the question of the homogeneity of physical properties within

subclasses and the variation from one subclass to

another.

I again stress that we have had no detailed

discussion of the taxonomy of guasi-WR objects, nor have we discussed their spectra or structure. This

represents a glaring neglect of what could be a powerful approach to the understanding of WR stars.

I am therefore putting quasi-WR objects at the head

of the list for today's discussion.

2. Specification of Quasi-WR Objects

I would recommend as candidates all stars with

broad emission lines and all stars with some pattern

similar to that of the two WR sequences. Thus we should certainly include P Cygni stars, some stages

of novae, nuclei of planetary nebulae, and possibly

the rocket-UV spectrum of the Sun and, by inference,

of many other stars.

3. Spectral Characteristics of the WR Atmosphere

This question was to be covered from two com- plementary aspects: Kuhi was to sximmarize the empirical information from which we might establish relationships between spectral and atmospheric features;

Anne was to summarize the diagnostic attempts at

interpreting spectral lines in tems of atmospheric

structure; from this we might possibly have expected

to deduce something about the conditions under which

the spectrum is produced.

To me, three things stand out from Kuhi ' s talk

because they seem to be important boundary conditions on any inference we may make about the state

of the atmosphere.

First, the breadth of the spectral emission

245
lines still seems inescapably linked to a velocity

dispersion of some sort. In the naive, thin-atmosphere interpretation, the existence of two definite

kinds of spectral lines - flat-topped profiles, often accompanied by violet absorption edges, and

more-or-less Gaussian profiles - is evidence for

both expansion and very high velocity random motions.

And indeed this is precisely the basis for Kuhi '

schematic model of the atmosphere: a photosphere

surmounted by a "turbulent" region surmounted in

turn by an expanding region.

Second, our one piece of evidence for stratification, Beals' old correlation between line width

and ionization stage, is still strongly supported

by Kuhi

But third, the binary character of many of the

best observed systems seems to play havoc with the

stratification problem. The behavior of lines during

eclipses emphasizes the point of Hanbury Brown's

question: Is it possible that an atmosphere which

extends to about 1/3 the separation of the components can follow the star around in its orbit? Can

we discuss it in terms of model-atmosphere computations for homogeneous static atmospheres?

From this it would appear that we have no clearcut, unambiguous information on the distribution of material, excitation, or electron temperature in the atmosphere. Indeed, I am sure you will

have noticed how careful Kuhi has been throughout

the discussions to avoid committing himself on any

of these subjects. But I personally think we have

a great deal of information, although clearly not

as much as we would like. First, no one seems to

dispute that Tg 5x10 '^^K in the region where the

continuum is formed nor that lines of N V and 0 VI

appear in at least some spectra- Since these lines

cannot be formed at 5x10 '*°K, Te must increase outward, at least initially. Second, we have data on V444 Cyg from the analysis by Mrs. Shapley and Kopal

,

which gives a rough distribution of density with

height over part of the atmosphere. As I showed

long ago, this puts stringent conditions on the

atmospheric model. Third, we have Kuhi ' s suggestion

that those lines with violet absorption edges are formed farther out in the atmosphere than the other

emission lines. He bases this on (1) the fact that

these absorption components are sharper than any

emission line and may therefore be formed outside

the electron scattering part of the atmosphere, and

(2) the fact that these absorption components arise

from levels strengthened by dilution effects, which

246
increase with increasing distance from the star. I think this suggestion deserves considerable attention since it could give important information on

the distribution of excitation and on the region of

expansion. Finally, the rocket-UV observations allow

us to explore the distribution in depth of excitation and ionization. As in similar solar investigations, this may at least provide limits on the

kinds of models that are compatible with the data.

Turning to the diagnostic aspects and the

summary by Anne Underhill, I find a slight difference of opinion between Anne and myself. Anne believes that everything is now too mixed up to con- clude much of anything. I, on the other hand, believe that we have significantly narrowed the alternatives for the atmospheric structure and also for

the causes of this structure. Let me quickly sum- marize the arguments I gave during the discussion

of her paper.

1. The spectrum consists mainly of emission

lines; absorption lines occur only as satellites

at the violet absorption edges of the emission.

This is a stringent condition. Of the three ways

to produce an intrinsic emission line, I have re- jected two (a Schuster-type mechanism and a fluorescent mechanism) , leaving only the possibility

of a chromosphere-corona type atmosphere in the

whole region of line formation. The alternative to an intrinsic emission line is one arising wholly

from the geometry of an extended atmosphere.

2. The fact that the excitation level of the

line spectrum exceeds that of the continuum ensures

that the lines are not formed, as in planetary nebulae, by a recombination spectrum. The situation

resembles that of a chromosphere-corona configuration.

3. The breadth and shape of the line profiles

also support the chromosphere-corona model. First,

the observed density gradients (V444 Cygni) and

atmospheric extensions imply either a dynamic support of the atmosphere or a temperature of about

10 ^°K. Random motions of the type suggested by the

profiles ensure mechanical dissipation of energy

and hence an energy supply. So while the emission

lines could conceivably be explained wholly by the

geometry of an extended atmosphere, the extended

atmosphere itself requires just that chromospherecorona configuration that it was introduced to obviate.

I conclude that the WR atmospheric configuration is simply that of a greatly enhanced chromo247
sphere-corona: enhanced in terms of the size of the

mechanical energy and momentum effects and in the

amount of mass involved in these effects. Given

this, we can comment on question (2) as well as on question (3) . The feature of the WR object that

produces its distinctive atmosphere is a high ve- locity field of mechanical energy; the features we should look for in the spectra of quasi-WR objects

are just those that accompany such a supply of

mechanical energy.

To clarify our picture of the WR and quasi-WR

objects, we will want to study:

1. The aerodynamics of motions of 100 to 1000

km/sec in a gas whose density is <. about 10 particles/cm^. In particular, we will want to inquire

into the diagnostic separation, in optically thin

and thick atmospheres, of random motions from motions

with a unidirectional velocity gradient along a radius.

2. Problems of radiative transfer under such

conditions

3. Problems of statistical equilibrium under

such conditions.

This now brings us to Lindsey Smith's remarks

on question (4)

4. The Structure of a WR Object and its Effect

on the Atmospheric Features

We must distinguish carefully between Lindsey

Smith's specific suggestion for a possible structure,

and the general kind of structure implied by the

physical nature of her suggestions. The specific

suggestions are a blend of the work by Paczynski and

his collaborators, and by Kippenhahn and his collaborators, as applied by Lindsey to support her inferences on the significance of spectral subclasses.

They also follow from the observation that a great

many of the WR stars - all of them, according to some conjectures - are components of close binary

systems, of which the other member is an early type

star usually of class 0.

The arguments rest on three properties of the

atmospheric model: (1) the requirement of a source

of mechanical energy; (2) the division into C and N

sequences, which can plausibly be interpreted as a difference in chemical composition; and (3) the

difference in ages, luminosities, and levels of

excitation between the subclasses and between the

sequences, coupled with evidence that the WR stars

248
are not very old objects. The arguments can be

summarized as follows: Because the primary component of a close binary loses mass to its companion,

it will evolve rapidly to a configuration with a thin outer hydrogen shell and a He-burning core.

It will then be overluminous for its mass. Such a helium configuration is probably vibrationally un- stable and if so, this instability could provide the

required supply of mechanical energy. The only

question is the extent to which the hydrogen shell

will damp the instability. The division into two

sequences is explained by a suggestion of Iben ' s

,

quoted by Paczynski, that an inner region where C

has been changed into N by the CNO cycle will be

separated by a sharp boundary from an outer part

where C is essentially unaffected. If the mass loss comes from sufficiently deep regions, then it could

change the C:N ratio. Paczynski gives other arguments {Acta Astronomioaj 1_7, 375-6 , 1967), but this

summary suffices to exhibit his thinking.

Finally, Lindsey acknowledged that there might

be other evolutionary paths by which a single star,

for example, could reach the He-burning configuration, but that these would take a longer time and

would result in a greater age for the star. This

agrees well with her results suggesting that single

WR stars may be slightly older than binaries.

Because the binary path seems well established,

our primary attention is directed toward these alternative evolutionary paths. We ask whether such

paths will lead to a WR object or only to a quasiWR object. What are the basic characteristics of

the model?

Fundamental to the model is the production of

a He- or C-burning core and thin H-rich shell. It

is this configuration which produces the instability which in turn produces the oscillations that

provide the mechanical heating of the atmosphere.

The characteristic of the WR spectrum that distinguishes it from, for example, the spectrum of the

outer solar atmosphere is the intensity of the effect. The entire line spectrum of the atmosphere

is involved. So it remains to be determined whether it is the quality or just the quantity of the

mechanical heating that is distinctive. But if a He- or C-burning core surmounted by a thin H envelope is a sufficient condition for a WR spectrum,

then anything that produces it will give the spectrum. For example, if a nova were a star which had

a He- or C-burning core, and if the nova explosion

could blow off enough of the outer H envelope, then

249
we might expect a WR spectriom. Indeed, any star

with such a core-configuration that could blow off

the H envelope should act the same way. Then to

what extent does the spectrum depend upon the re- maining characteristics of the star, such as mass,

radius, and gravity? The details of the energy

supply and the accompanying aerodynamic motions are

at this point mere conjecture.

This is a good point at which to close a sur- vey of questions (3) and (4) . In discussing the

atmosphere we have given reasons for rejecting some alternatives as well as reasons for preferring others.

Their consequences have not been worked out, but it seems to me that we have only a small range in possible atmospheric models. The discussion of a general model for WR objects is provocative, even though highly speculative. But this is what we need

for the WR objects: something on which to hang our hat as a first approximation.

5. The Structure of Quasi-WR Objects and Its

Effect on the Atmospheric Features

We have had no detailed discussion of question

(5) . I have commented over the years on the relation between chromospheres and normal atmospheres

in various kinds of objects, but I shall not repeat

them here. It is clear that there are many sources

of stellar instability and that they all lead to the

production of chromospheres and coronas of one kind

or another. Whether the variations on the He- or C-burning core and on the H-rich envelope lead to

the unique type of chromosphere-corona that produces

the quasi-WR spectrum can only be settled by further

investigation. In the meantime, the solar and WR

features appear to represent two extremes in the

range of mechanisms by which such outer atmospheres

can be produced.

These points are now open for general discussion.

250
DISCUSSION

Thomas : 1 suggest we follow the outline of the

siammary, treating in succession questions (1) to

(5) . So we start with (1) : taxonomy. We discuss

the possible existence of a homogeneous class of WR

objects or phenomena exhibiting a "pure" WR spectrum. We also consider the quasi~WR objects. The

discussion should be primarily from an observational-empirical viewpoint; Wc should recognize, however, that we are likely to become involved with

considerations on question (2) - those features com- mon to both WR and quasi-WR spectra that imply that

some atmospheric properties of the former are also

present in the latter.

Underhill: 1 don't really see the point of

separating the WR spectra into these two categories,

WR and quasi-WR. In all these types of stars we are looking chiefly at emission lines. The excitation

of emission lines is a spectroscopic phenomenon that occurs in a plasma of a certain density and a certain

amount of excitation energy. We have concluded that

we can not maintain this level of excitation by ra- diation alone; we have to have a source of mechanical

energy. So from the physical viewpoint, all these

objects belong together. The so-called classical WR

stars are distinguished only by their distinctive

and well grouped class of lines. The others show

the same phenomena, but because of the level of ex- citation or something - it's hard to be specific when

you haven't described the physical process - we are not so struck by their peculiarity.

Thomas: Precisely. The "pure" WR spectrum is

well defined by a set of characteristics which are exceptional and which I believe require very precise

conditions for their formation. The quasi-WR spectra have no such well-defined characteristics, nor

do they imply, in consequence, such a unique set of

atmospheric conditions. I think you are much too

glib in passing off a spectrum composed exclusively

of emission lines as "a spectroscopic phenomenon

that occurs in a plasma of a certain density and

level of excitation". Can you really tell us the

density and excitation conditions required to produce the WR emission lines? I remind you of the

long-standing controversy in the astronomical literature on the interpretation of the conditions underlying the production of any emission line. Only

251
recently have we even begun to sort out the situations under which such-and-such a variety of emission line can be produced.

Underhill : The difference between your two

general classes is not well defined. The various

spectral types are well defined and easily identified, but the basic physical problems are the same

for all of them. So even though we speak of 7 or

8 different spectral classes, we should realize

that all these stars lie in one physical class - they are stars with spectroscopic phenomena that

imply a mechanical source of heating in an extended

atmosphere

Thomas: Welcome to the ranks of those who

support the broad general thesis of twenty years

ago. But now we are trying to establish the details

of the mechanical energy supply and the precise way

in which it interacts with the atmosphere. On the

basis of what we have learned about the interpretation of spectra produced by such atmospheres, this

division seems helpful in establishing reference

points for our analysis.

Hjellming: I think the two general classes

are based more on our ignorance than on our knowledge of any given object. First we define a WR

phenomenon; then if we know nothing more about the

object than that it shows a WR spectrum, we call it

a WR object; if we know something more about it - for example, that it has been a nova - we put it in

another category.

Stecher: I would like to correct an omission

in the category of quasi-WR phenomena. Every observation that Carruthers , Morton and I have made

on supergiants earlier than B8 (and B8 is the latest

type we have observed) shows the resonance lines in

emission with a Doppler-shifted absorption edge.

The velocities corresponding to this shift are of

the order of 2000 km/sec or greater, twice that observed in WR stars. I interpret this as evidence

that mass is leaving early type supergiants with

very large velocities.

Thomas : Does this mean that you would include

all supergiants earlier than B8 in the quasi-WR category?

Steoher: I think we have to include them. Except for the presence of forbidden lines and recom- bination lines, they look very similar to the WR

stars in the rocket UV. Perhaps the higher velocity

fields give lower densities, so that you can't populate the higher levels that give some of the WR lines.

Roman: I would like to defend the distinction

252
between WR objects and quasi-WR objects. I don't

know of any system of spectral classification in

which you do not have both borderline cases and re- lated objects that do not satisfy any firm set of

criteria. That doesn't make the division of objects

into spectral classes useless. I think you'll have

to keep your eye on quasi-WR phenomena as an indication of what is going on in WR objects. But I think

that if you lump them all together, you will become

hopelessly confused in trying to define a WR object.

Sahade : I too agree with this separation into

two groups. I would just like to add a fifth condition to the four defining group A: the time-constancy of the four other conditions.

With regard to Of stars, which have been included in group B, I think we should keep in mind

that they might turn out to be very similar to the

WR objects. Let me remind you of the characteristics they share in common. Those Of stars that are binaries seem to have 0-type companions and seem to

be less massive than their companions. You have already seen the similarity in spectral features:

Stecher showed it for C Pup in the rocket UV; several

years ago R. Wilson of Edinburgh showed that underlying the narrow emission features characteristic of

spectra in the photographic region, are broad emission features similar to those found in WR stars.

So it may well be that the Of stars are closely re- lated to the WR stars in some way not yet clear to us. West: I suggest we change the name "WR phenomenon" to "emission line phenomenon" because all

these objects are special cases of the same mechanism.

Thomas: This would be difficult: after all the

Sun is an emission line object, so you are probably

suggesting we include most of the stars. A more

specific objection is that many kinds of nebulae

are emission line objects, and we do not believe,

for example, that the emission lines in planetary

nebulae are produced by the dissipation of mechanical energy, but rather by the reverse process, a degradation of an ultraviolet radiation field.

Steoher: The question is: What is the natural

division? We have been discussing the spectral

lines as the distinctive feature; but I think it may

be the velocity field which provides the physical

relationship between these objects. I think Miss

Underhill has been implicitly assuming a large mass

loss from a number of these objects, and this is

what relates them physically.

Thomas: Velocity fields and mass loss are

253
quantities derived from an analysis of spectral

lines. Until you are sure your diagnostic methods

are correct you cannot use them as taxonomic criteria. And a large part of our discussion has centered on the methods by which to analyze these spectra.

I suggest we now move on to discuss the homogeneity and uniqueness of these spectral subclasses

with respect to the physical quantities characterizing WR stars. Each of these quantities must be derived by some kind of analysis of an observed quantity.

Smith: I have always thought the things that

define what a star is and what it looks like are its

basic properties: mass, initial composition, and

age. Now in the present discussion, we may have to

add as an additional parameter the binary character

of the star, and therefore the properties of the com- panion and the separation may be important. Given

all these properties, we should be able to predict

the evolution and spectral characteristics of the

star. So the simplest explanation of these objects

would lie in a one-to-one relationship between the

spectral properties we observe now and the initial

parameters. If you claim that a great range of initial parameters can converge to produce one spectral

subclass, the situation becomes extraordinarily com- plicated, and I then see no reason why we should

have a large range of subclasses. You would be saying the WR subclass is some sort of random property

accidentally generated from one of a series of initial parameters. It seems infinitely more likely,

and is in agreement with observed correlations , that

certain initial parameters will produce certain final products. So I would suppose that when we see a class that appears homogeneous in spectral properties, luminosity and distribution, this class does

represent a stage in a unique class of objects (as

defined by the initial parameters) . Therefore I as- sert that if we define our subclasses properly,

those properties I listed will lie within narrow

limits in each subclass.

Thomas: Do I understand properly? You assert

that luminosity, radius, mass, and some parameters

characterizing the continuous radiation field and

the binary character should be nearly constant within a given subclass?

Smith: Yes, within a reasonable standard deviation, which represents not just observational

error but a genuine spread in physical properties.

And some of these properties vary from one subclass

to another.

254
Payne-Gaposohkin : I don't think I can let

this point of view pass. There are a number of parameters mentioned here, but the one I shall talk

about is mass. It would be a brave man who would

make the statement that throughout the whole range

of spectral classes there is even one class that has

a unique mass. There are less than 60 stellar masses known, including both components of binaries, and

you must have a binary to get a mass. Those binaries from which masses are determined can be roughly

classified as close binaries and wide binaries. I don't think anyone will quarrel with the idea that

the wide binaries such as ZZ Boo, WZ Oph, 70 Oph

are best for determining masses. Now these systems

all have main sequence components, and if there is

any spectral class for which you can state that you

can unequivocally assign a mass, it would be mainsequence luminosity class V. Here I wouldn't quarrel with Lindsey.

However a large number of masses are not determined from wide binaries on the main sequence. Those

stars which have the very highest masses, e.g., V Pup,

y Scorp, UW Can Maj , are almost contact binaries.

They are still main-sequence stars, but in terms of mass exchange, they are somewhat contaminated. Nonetheless, these stars do not contradict Lindsey 's as- sertions. If we include luminosity class as well as spectral class, stars of similar spectra do seem to

have similar masses. You must recognize, however,

that the material is rather more uncertain because

of uncertainty in the inclination of the orbit.

Now there are other stars, still on the main

sequence and still contact binaries, for which the masses are notoriously different although the spectral classes are nearly the same. The W Ursa Majoris class contains a niamber of examples, e.g., W Urs

Maj and U Peg. Then you have another set of still

queerer pairs for which the masses are extremely

well determined: Sirius and its companion, and

Procyon and its companion; Z Here consists of 2 subgiants of about the same mass but different spectral

class and too bright for the mass-luminosity relation. There are still odder objects like 85 Peg,

which has two components of apparently equal mass,

one a main-sequence star, the other 3 magnitudes

fainter and probably a subdwarf . All these stars

form strong exceptions to Lindsey 's suggestion of

a specific mass for each spectral class. There is

also a very large group of stars off the main se- quence, £ Aurig , C Aurig and W Ceph , for which

masses are known. Both components of C Aurig seem

255
to satisfy the mass- luminosity relation, but some

of the others don't.

All this simply illustrates Strand's statement

that stellar evolution is all very well until you

begin to look at double stars. It is nice to think

you can get out of the problem by saying all the

problem children are contact binaries; but you can't

use that escape for Sirius or for Procyon, which are wide visual binaries with periods of about 50 years.

It is clear that when you speak of stars" with masses

that do not agree with the mass-luminosity relation,

you are speaking in the majority of cases of close

and contact binaries. And I think it is in this

category that we must place the WR stars. If I might stretch a point, I think we can regard WR stars

as resembling Algol stars in that the two components

are physically so extremely different; one is on the

main sequence, and the other is probably a subgiant.

How many well-determined masses are there among the

WR stars?

Underhill: Seven systems - possibly eight - of

which none are well determined.

Smith: Two and one-half.

S» Gaposohkin : Only one star.

Payne-Gaposchkin : I am talking about M, not

M sin i; if they are not eclipsing systems, you can only determine M sin^ i.

Underhill: There are 3 eclipsing systems, not

counting CQ Ceph.

Thomas: Cecilia, it seems to me you are saying: (1) We only know one or two masses, so it is

impossible to know whether Lindsey's assertion is

correct, but (2) on the basis of my other experience, I, C.P.G. , do not believe it.

Payne-Gaposchkin : From what we know about far more well-behaved stars than the WR's.

Thomas : Would you comment on the visual magnitudes from the same standpoint? You are destroying our picture of physical homogeneity within subclasses; I am asking you to destroy it point by point.

Payne-Gaposchkin : Without having the data before me, I would not like to comment. I have heard

what Lindsey and Anne have said , but I haven't been

looking at their data. The absolute visual m.agnitude depends critically on the WR stars in the Magellanic Clouds. But turning again to stars in general, I think that except on the main sequence it

would be very dangerous to try to specify the absolute magnitude from the spectral class. This is

particularly true for the supergiants.

Thomas : We find ourselves in a very interest256
ing situation. My viewpoint is similar to Lindsey's.

If we can't get homogeneity by progressing to smaller

and smaller subclasses, we must discuss each star

individually. This is grim. So let's consider the

next point: the possibility that a temperature, or

some other measure of the distribution of energy in

the continuum is constant within a subclass.

Payne-Gaposchkin : Here I don't object. What

evidence we have is sound and convincing. But there

are mighty few stars , and you can always put a straight line through two points.

Underhill : 1 want to make a basic comment on the method of procedure. For absorption-line stars on the main sequence, we have some confidence in the

idea that if the spectrum is excited by radiation,

the mass and luminosity are related to certain identifiable spectral features. We have hypothesized

that this is a fixed rule of nature, and that we know what the rule is. I am not so certain that we

do know the rule, but perhaps we are close to it.

Payne-Gapos chkin : So you disagree with my

statement that knowledge of the spectral class does

not necessarily imply knowledge of the mass?

Underhill : That is the assumption on which

we proceed. I agree with you that it is a bit

doubtful, but it is the basic assumption that is

made - for the main sequence and nowhere else. Now

it will certainly not work for the WR stars or for

the whole group of stars with emission lines. There

we are still searching for the relationship between

spectrum and physical characteristics such as total

radiation field and mechanical energy. When we can isolate the basic physical processes, we shall be

able to relate them to the one or two masses and one

or two luminosities that we know.

Payne-Gaposohkin : What I really want to take

issue with is the belief that because you have to

know certain parameters in order to interpret the

spectra in terms of a given theory, it is correct

to assume that you do know them.

Thomas: Let us carefully distinguish two questions: The first is whether in principle and in

practice there _is a homogeneous set of physical

parameters that characterize a subclass; the second

is whether we actually know the values of these parameters for each subclass. I had thought you were

saying "no" to the first question; your last comment

seems to imply "possibly" to the first question but

at the moment "no" to the second.

Payne-Gaposohkin: I don't think that even in

principle there is a homogeneous set of parameters

257
for non-main sequence stars. I put most of my em- phasis on masses only because they are a definite

thing to talk about « I think Nancy Roman knows more about masses than I do, and I wish she would say

whether she agrees with me.

S, Gaposchkin : There are only 54 masses determined from eclipsing binaries; if we include the

visual binaries, the figure rises to no more than

70. The majority of these lie on the main sequence.

Underhill : Everything we know about the visual

absolute magnitudes of WR stars is contained in the

following table, extracted from Lindsey Smith's

thesis. The results are b;="^e.d on stars in the Large

Magellanic Cloud that show no evidence of being double. You can see how few stars we have; you can Class Number M

V

S.D.

WN3 2 -4.5 ±0.1

WN4 5 -3.9 ±0.3

WN5 2 -4.1 ±0.05

WN7 4 -6.8 ±1.0

WN8 3 -6.2 ±0.4

guess whether a statistician would go along with

conclusions drawn from them. In addition, we have

Y Vel and one other WC8 in the Galaxy, which Lindsey

suggests gives us a magnitude about -6.2. I have

stated why I think -5 would be more appropriate.

So if I put this all together, I have WC7 and WC8

about -6 mag or brighter and all the rest about

-4.5 ± whatever your statistical sense tells you is

appropriate. I hope you will agree that this material is not sufficient to permit any differentiation with respect to spectral subclasses.

Thomas: The question was not whether a given

mass or luminosity represents a unique subclass,

but whether a given subclass represents a unique

mass and luminosity - within some range.

Underhill : You can't answer either question.

The absolute visual magnitudes are in no way related

to mass. For the WR stars, I don't need to know the

subclass - the absolute visual magnitude is -4.5.

Thomas: Anne! If I know the spectral subclass,

do I know the visual magnitude?

Underhill : Yes. Even if you don't know the

subclass, you know the visual magnitude.

258
Payne - Gap o s ohkin : A standard deviation of ±

magnitude means you don't know!

S. Gaposchkin : I have the best determined

absolute luminosity: V444 Cyg (WN7) , -2.7. How can you disagree with that?

Underhill : The light curve is ambiguous. I know how to interpret the light curve of two spherical, well-separated stars. I don't know how to

interpret the light curve of V4 44 Cyg. All spectroscopic investigations of that system indicate that

these stars are not well separated and they are not

spherical. Therefore I do not believe your value.

Kuhi : I agree with Anne one hundred percent.

Sahade: The point is that the luminosity ratio was not determined from the light curve because it was impossible. It was determined from spectra of

stars other than V444 Cyg. Kron worked out his

light curve and had to use the luminosity ratio that was in the literature, but that value was completely

unreliable.

Westevlund : I would like to offer the following explanation for the lack of agreement between the

absolute magnitudes determined from binaries and from

WR stars in the Magellanic Clouds. In the Magellanic

Clouds, the Wolf-Rayet stars appear at the turn-off

points of the HR diagrams of clusters and associations. The ages of these stars are about 5x10^ years.

During this time the massive stars have increased in

luminosity by at least 1.5 magnitudes. From the binaries you derive a luminosity by assuming the star

is still on the main sequence. This value must then

turn out to be at least 1.5 magnitudes too low. I think everyone will agree to that. The WR stars are not much more over-luminous than other slightly

evolved massive stars.

Fayne -Gap OS ohkin : On homogeneity of subclasses,

I think it is worthwhile to raise again the question

of the binary character of WR stars. If one looks

in the catalogues, one sees a large number of doubleline spectroscopic binaries. There is also a large

number of single-line spectroscopic binaries in which

the second star is there, but is not bright enough

to affect the integrated spectrum. Its presence can be detected only by a change in radial velocity that

can't be ascribed to pulsations of a single star. If

there are no absorption lines, it is very difficult

to detect velocity variations in WR stars. I don't

remember the propoiTtion of double- to single-line

spectroscopic binaries: Is it not possible that

double-line spectroscopic binaries are in a minority

and that many WR stars are members of single-line

259
spectroscopic binary systems? This would make a big

difference in all our discussions about WR objects.

Underhill : One-third of all WR stars might be

members of spectroscopic binary systems; we don't

know about the others.

S. Gaposchkin : I too believe that all WR stars

may be members of binary systems

Thomas: Now today we are trying to summarize

areas that were discussed in preceding sessions and

to raise new considerations that may have come up

when you were digesting earlier discussions. If

there are points on which we agree, it will be sur- prising but most welcome. If there are points on which we disagree, we should carry the discussion

just far enough to focus on the issues and then drop

it. Hopefully you will go home, solve the problem,

publish a paper, and acknowledge this symposium for

providing the choler that pushed you into doing it.

Now some of you are outspoken in your disagreement;

some of you don't come forth with your objections

except during the coffee break; I have been asked to

clarify two points raised over coffee.

The first point is Charlotte Pecker-Wimel ' s as- sertion that my criteria (1) to (4) , by which I define the WR spectroscopic class, simply confuse the

issue of the central stars of planetary nebulae.

She asserts that the planetary nuclei with WR spectra satisfy all the criteria (1) to (4) and that the

only way to distinguish them observationally is by

adding another criterion specifying the range in absolute magnitude. But absolute magnitude is not an observational criterion. It is my understanding

that our values for the luminosities of these stars

depend on distance determinations by Seaton and others who warn that they are subject to large errors. Similarly, we believe that planetary nuclei have masses of about one solar mass, but again this is

not an observational criterion. Now you will recall

Sahade's suggestion that as a fifth criterion we specify the time-independence of criteria (1) to

(4) . If someone can add a further specific criterion, I shall be delighted. But I would have thought

we had just had an excellent demonstration of the

uncertainty in the absolute magnitudes of WR stars.

Aller's discussion of his own and Lindsey's work on the nuclei of planetary nebulae covers about all we

can say.

The second point raised over coffee concerns

the definition of the chromosphere-corona. By this

I mean that part of the atmosphere where the assumption of radiative equilibrium breaks down, where a 260
mechanical supply of energy is required to maintain

the state of the atmosphere. In the classical model

of a stellar atmosphere, the temperature decreases

monotonically outward; thus the chromosphere-corona

is often thought of as that part of the atmosphere

lying above a temperature minimiam. But several

years ago Cayrel suggested that the same mechanism

that fixes the temperature of planetary nebulae at

about 10'*°K would act to produce an outward increase

in temperature even in a classical model; would

rise from the boundary temperature to about the effective temperature of the star. The explanation is

that at low densities Tg is fixed by the spectral

quality of the radiation rather than by the energy

density of the radiation field as at high densities.

So if we want it to connote a departure from the

classical model, we have to define a chromospherecorona in terms of the contribution from non-radiative (usually mechanical) energy sources. Johnson: So your definition of a chromospherecorona could include an envelope produced and heated

by mass exchange from one component of a binary to

the other?

Thomas: Yes. Now that we have discussed the

question of WR objects and the homogeneity of their

subclasses, I suggest we turn to question (2) : Are

there features common to both WR and quasi-WR spectra that are produced by the same mechanism in each

case?

Sahade: It is not obvious that the same mechanism produces the instability in all your quasi-WR

objects. And I would have strong doubts about including the Be stars as members because their instability seems to be gravitational, and this is ruled

out for WR stars.

Thomas: I think you must distinguish carefully

between the mechanism of the instability and the

physical effect of that instability on the atmosphere

or its spectral features. It seems to me we are now pretty well agreed that those objects exhibiting some variety of WR spectrum have in common a supply of

mechanical energy to their atmospheres. In proposing my categories A and B, I hoped to suggest differences in the detailed properties of this energy

supply. We might suppose that the classical WR

objects have a common kind of instability which produces the particular size and form of mechanical energy supply responsible for the "pure" WR spectrum.

The quasi-WR objects might have different kinds of

instabilities, which lead to different properties of'

the atmospheric energy supply and thus to different

261
spectral features. If the common spectral features

imply common atmospheric conditions, their study may

help to clarify the mechanisms which produce them.

I think the chief criterion for a quasi-WR class is

a functional one: how much will it contribute to our study of classical WR stars? We could, for example, try to limit the category to those objects

that have a large momentum as well as a large energy

supply to their atmospheres; but then we would be

prejudging on a basis not wholly observational.

Peokev-Wimel : You must be cautious about including the central stars of planetary nebulae in

your list of quasi-WR objects because some of them

show only absorption lines, and some show only continuum.

Roman: I'd like to be sure that "novae" include

old novae and dwarf novae as well as supernovae. And

I think you should include symbiotic stars and possibly those stars that have been definitely identified as x-ray sources. Payne-Gaposohkin : 1 would like to draw your

attention to the similarities between novae and

Wolf-Rayet spectra. First, they both have broad

emission lines with violet-displaced components.

There is a definite correlation between the violet

displacement of the lines, presumably due to the

ejection velocity, and the apparent spectral class

of the star. Nova Pictoris, one of the slowest

novae, had an expansion velocity of 75 km/sec and a spectrum of class F5 before maximum. Nova Persei,

1901, was one of the fastest with an expansion ve- locity of about 2700 km/sec and a B-spectrum before

maximum. Various other novae fell in between. Also,

the excitation of the spectrum associated with the

absorption components increases as velocity increases;

this convinces me that it is largely mechanical. My

second point is that among novae, as among Wolf-Rayet

stars, there appears to be a significant difference

in chemical composition as determined from the bright

lines. For example. Nova Persei showed the forbidden

lines of Ne III, Ne IV, and Ne V in successive stages.

Nova Aquilae, on the other hand, was strong in oxygen,

especially 0 III. Nova Geminorum 1912 was very strong

in nitrogen. Nova Pictoris was strong in the iron

lines, all the way up to Fe VII, and oxygen lines were very inconspicuous. Does this really reflect a difference in chemical composition? I know of only

one nova for which a curve-of-growth analysis has

been attempted. This is Nova Herculis, for which I made the analysis myself, but the results cannot be

taken very seriously. However, there does seem to

262
be a high abundance of carbon (recollect the cyanogen spectrum of Nova Herculis) and also of oxygen.

I believe Aller attempted a similar analysis which

has remained unpublished. Thus the situation for novae is similar to that for WR stars: we have neon

novae, oxygen novae, nitrogen novae, and iron novae. The problem is, however, more difficult for novae

than for WR stars because the most conspicuous lines

involve forbidden transitions. I would guess that

the WR problem will be solved before the nova problem.

For these two reasons - the correlation of ex- citation with radial velocity and the apparent difference in chemical composition - I feel we are justified in including novae in the same general category as the Wolf-Rayet stars. The novae certainly

have expanding outer regions; for many of them we

can directly observe the expanding nebulosities.

Undevhill : I think the novae spectrum gives

every indication that it is excited by mechanical

energy. In general the novae have lower expansion

velocities than do the WR stars, although there is some overlap. It seems to me probable that the ex- citation of the spectrum is a local process, depending upon the density. If you talk to plasma physicists who try to excite the third, fourth, and fifth

spectra of various ions in theta-pinch configurations,

they tell you that the geometry and local density are critical in determining which spectra they get. So

I am more than ever convinced that the apparent differences in WR spectra and in novae spectra are largely chance effects of density and collisional mechanisms. Until we have sorted out these things, I am dead against speaking of abundance differences.

Thomas: I would agree with almost everything

you say, except when you refer to "chance occurrences" and to the spectrum being excited locally.

The latter is a strong statement; it implies that

the radiation field produced in the atmosphere has

nothing to do with the excitation; it implies that

the atmosphere is optically thin. Maybe so, but it

has yet to be- proved.

Stecher: In trying to decide between protons,

electrons, and a particles as the exciting particles

in collisions, we should note that because of con- servation of spin, neither protons nor a particles

can easily excite the triplets in C III, N IV and

O V. Recent experiments have shown that the cross section for exciting the He triplets is down by a factor of at least 100 over that for exciting the

singlets. Thus although you can ionize in this way.

263
the spin conditions make it unlikely you will ex- cite the upper triplet levels.

Underhill : A more important point is that

protons and electrons are equally effective in ex- citing He when they have the same velocity, not the same energy. So their effectiveness depends on how

they are being accelerated. If we have magnetic

fields, the proton excitation may be important; if we have only thermal velocity fields, I agree that

the electrons are the most important.

Kuhi : I'd like to add one point, to emphasize

Anne's concern. Mrs. Gaposchkin mentioned Nova

Aquilae; this nova has changed its spectrum quite

dramatically over the years. It started with very

strong N III lines, the A4640 complex, and then

gradually changed over to C III and IV, X4650. This

is a very important point because surely the composition didn't change.

Beats: Rather than your two categories of

classical -WR stars and quasi-WR stars, I think that

I would prefer the following breakdown: (a) regular

non-binary WR stars; (b) regular binary WR stars;

(c) WR nuclei of planetary nebulae; (d) stars showing only He and Ha. It appears that group (a) is

getting smaller all the time, but it is the one I would study were I starting out again. Group (c) seems to be a good deal smaller than groups (a) and

(b) , but this is hardly a reason for excluding it.

I had never heard of group (d) before, but it looks

highly interesting. In a print Miss Smith showed

me. Ha is so strong and He II X5411 so weak that I don't feel one could attribute it entirely to He.

I think we must consider class (d) , particularly in

view of Miss Smith's discussion of the evolution of

these stars.

Now I have two remarks that may possibly re- late to the question of spectral features common to

WR stars and other objects. One concerns the suggestion of a coronal heating mechanism, which I find very interesting; the second is another solar

analogy related to flares. I understand that in some solar flares, atoms are ejected with velocities

up to 1200 km/sec. Furthermore, flares are associated with solar prominences. So if these kinds of

things can happen on a quiet star like the Sun, they

might occur in a greatly magnified form on the WR

stars. In complicated conditions like those responsible for WR emission, it is unlikely that only one mechanism is operating. I would suggest that the

coronal and the flare mechanisms should be considered.

264
Ej ellming : Following in the spirit of Beals'

remark, I would like to raise the possibility of a classification scheme based on the possible forms

of mechanical energy. From the solar data, we can make two very clear divisions. First, there are the

hydrodynamical phenomena such as the corona and the

solar wind, and second there are the more energetic

phenomena related to cosmic rays. It is well known

that there are roughly three classes of cosmic rays

from the Sun. One class occurs in the presence of a strong flare and produces extremely energetic particles up into the bev range; this happens once every

year or two. A second class, again associated with

a flare, produces energies of tens or hundreds of

mev and occurs once every few months. Finally there

is a quiescent, virtually continuous, form of emission which produces particles of a few mev. In discussing the forms of mechanical energy in the WR

stars, you could be concerned with either or both

of these classes. Everyone who has mentioned a non- Maxwellian velocity distribution could equally well

speak of a Maxwellian distribution of velocities

associated with hydrodynamical phenomena plus a non- Maxwellian tail which everyone has called cosmic

rays

Undevhill : All the more energetic • classes of

mechanical energy in the Sun have in common the presence of a magnetic field. So I think it is perhaps

serious that we have omitted mentioning that magnetic

fields are necessary for some types of mechanical

excitation.

Thomas: We don't seem to have made much headway with question (2), and in a sense Beals' remarks'

carry us over to question (3) : What are the distinguishing properties of the WR atmosphere that produce the unique WR spectrum? Kuhi has summarized our present knowledge of the WR spectrum, and Anne Underhill has commented on the diagnostics. I have men- tioned three points that seemed to me to stand out in

Kuhi ' s talk: the breadths of spectral lines, the

possible relation between line width and ionization,

and the complexity of the atmospheric situation as revealed by the study of binaries.

Sahade : As a bridge between this point and the

discussion of quasi-WR spectra, we should not forget

that novae and symbiotic stars may all be binaries,

and the mechanism at work in them may be connected

with their binary nature.

Deinzer: Kuhi suggested an atmospheric model

consisting of two regions: a "turbulent" region in

which the emission lines are supposed to be formed.

265
and an outer expanding region. Would anyone care to

make a guess at the mass of these regions? I am concerned with what happens to the mass which a WR

star loses during its lifetime. It must be somewhere

around a single star; in a binary system, it may go

to the companion.

Thomas: This was the problem underlying an

earli'er discussion of the nebulosity associated with

WR stars. The question was: Why, if indeed there is a general mass loss, do only some stars show nebulosity?

Lindsey's answer was that the nebulosity you see con- sists not only of the mass ejected by the WR star,

but chiefly of the matter swept up by this mass from

the interstellar medium. The ejected mass has so much energy and momentum that it sweeps up much more material than is ejected. You wouldn't then expect

to see nebulosity except in those regions of the Galaxy where there is plenty of interstellar material.

A detailed interpretation of the nebulosity will depend strongly on a knowledge of the ejected mass, on the properties of the interstellar medium, and on the

details of the interaction between the two. Also

note such arguments as Schmidt-Kaler ' s in a preceding

session, where he suggested that any ejection of matter could trigger the formation of a whole shell of

stars. So with regard to mass loss and its effect

on the environment, there is a great deal of work to

be done. Don't forget that not everyone accepts this

idea of mass loss from single stars; and even if you

do accept it, you still have to know how great it is

and how it produces the observable effects. This

brings us back to all the diagnostic problems discussed earlier: how do you infer the velocity

fields, the effect of electron scattering on line

profiles, and so on. Lindsey, do you have a figure

for the mass of the nebula?

Smith: My figures depend on all the things

you mentioned. The mass of the nebula lies between

1 and many hundreds of solar masses, depending on the

assumed mass loss, the density of the interstellar

medium, and the interaction.

S. Gaposohkin: There is no mass loss except

in the form of the so-called "jet" lines in 3 Lyr.

General chaos and screams of protest

Underhill : Certainly there is mass loss. It's

also easy to calculate the mass of the turbulent re- gion of the atmosphere. Just take the volume of the

shell contained between the photospheric radius of 5 solar radii and an outer radius of about a factor 7 larger; guess a density, and you have the mass. A

reasonable density is 10^^ particles/cm^. Assume the

266
shell is pure H, or the usual H:He mixture, and you

get about 10^^ gm.

Deinzer: And the mass in the expanding shell?

Underhill : Forget it - it's small. For an ex- pansion velocity of about 10^ km/sec, the upper limit

on the mass loss is about 10"^ solar masses per year.

Electron scattering may affect our estimate of the

expansion velocity, but not by an order of magnitude.

Deinzer: This means the WR stars can't be

very old.

Thomas: No, these estimates can be very bad;

the star could even have thrown off its mass before

entering the WR phase. Our interpretation of the

so-called expansion velocity could be very wrong,

considering the effects of electron scattering and

the difficulties of separating the expansion velocity from the velocity gradients. The picture of a turbulent lower atmosphere surrounded by an expanding outer atmosphere is highly qualitative, especially in view of the size of the velocities and the

lack of a physically consistent model relating them

to the excitation state of the atmosphere.

Hummer: I'd like to express a personal prejudice that electron scattering will not play a major role in determining the appearance of the

lines in a WR star. It would be difficult now to

support this with facts, but this is the impression

I get from having done a fair amount of work recently on electron scattering in line formation. An

important point to remember is that the electrons

may be in the same region as the atoms ^ and not on

top, in a reversing layer of the kind usually con- sidered in theoretical work on electron scattering

in line formation. Because of the competition from

absorption and line scattering, the effect of electron scattering is much weaker than if all of the

electrons were isolated in a reversing layer.

Johnson: Would you be willing to draw a typical line profile and show us where in the model

each part of the line is formed?

Hummer: I ' d be willing to speculate on the

kind of velocity fields associated with each kind

of profile. The characteristic very broad flattopped lines are formed by purely radial motion

and occultation, with not much optical depth effect,

probably because the velocity gradient is sufficiently large that even if the densities and path lengths

are large, everything is optically thin. The big

Gaussian lines are formed in turbulent regions in

which the turbulence de-saturates the line by spreading the opacity over a sufficiently large spectral

267
region that things can't become sufficiently optically thick for self-reversals to appear. The third

kind of profile is the P Cygni type, which comes from an optically thick region in which the velocity

in increasing outwards.

Wrubel: But you still can't explain Kuhi's

eclipse data with this interpretation.

Hummer: That's right. I've been talking about

simple geometries. If I understand what Len has

been saying, I would have to look at the radiative

transfer problem in a dumbbell.

Kuhi : Let's not worry about the eclipsing

system - the cases with simple geometries should be

solved first.

Nariai: May I ask if anyone here really thinks

the turbulent velocity is of the order 1000 km/sec?

It seems to me that this is what you are saying. I also recall your saying the temperature of the gas

which produces the line is about 3x10 '*'*K. Thomas: The empirical basis for the use of

"turbulence" is that the velocity distribution in

the atmosphere must have the same effect on the line

profile as a random distribution of velocities along

the line of sight over a volxame element whose size

is about the mean free path of a photon. Now whether

this effect can be mimicked by a distribution such

that the motions are all in one direction in one region of space and all in the opposite direction

in another region of space, or whether we require

random motions in a given region of space, depends

very much upon the size and distribution of the opacity, and the opacity in turn depends upon the velocity, so we need self-consistent models. For years

astronomers used the term "astronomical turbulence"

to indicate random motions in one region of space.

Collisions between turbulent elements were ignored

or were tacitly assumed to be absent or ineffective.

All these aspects of what one means by "turbulence"

in astronomy are vague. If one uses "turbulence"

in the aerodynamical sense of random motions in a given volume, then as I pointed out long ago this

implies that T^ 10 in the WR stars. I also

said I would prefer a more elaborate aerodynamic

explanation, because 10^ °K seemed high even to me,

but perhaps I was unduly pessimistic. In any event,

the 3x10** °K you quote refers to the region where the

continuum is formed. It is interesting to note that

the value of the electron temperature acceptable to

the astronomical public seems to be increasing with

time.

Hummer: The line opacity will be different in

268
a turbulent situation than in one with velocity

gradients, and therefore the relation of the source function to the emergent flux will be quite different in the two cases. Hopefully we can separate the

two effects by looking at the line profiles.

Steoher: 10 can be ruled out by the x-ray

data. The x-ray flux at the Earth from y Vel would

be stupendous if the entire shell were at 10 ^°K.

Thomas: You are probably right, but I warn you

that you must consider the mechanism by which the

x-rays are formed. I presume your remark is based

on the assumption of blackbody emission at 10 ^°K.

This means the shell must be opaque in the x-ray

region you are studying. Do you believe that? I remind you that in the early days when I argued for

a high-temperature solar chromosphere, the same argument was made with respect to the hydrogen-Lyman

continuum. When rocket observations indicated a radiation temperature of about 6x10 for the Lyman

continuum, the arguments were repeated. The point

is that these arguments are wrong; the non-LTE effects introduce a big correction factor; so we do in

fact predict the observed Lyman continuum from the

high temperature chromosphere, I caution you to

check the emission from your WR model on the basis

of a correct theory before you determine Te from

the x-ray flux.

Underhill : Don't forget that in these observations you are integrating over the whole disk. Now

suppose I make a schematic model with absolutely no physical basis. I postulate a magnetic field that

looks like a porcupine, and this produces coronal

streamers of the same geometry. By an unknown flare

mechanism similar to that in the Sun, streams of

particles with varying densities shoot out from each

of these streamers. As observed from the Earth,

this group of overgrown spicules has a wide distribution of velocities and will produce a Gaussianshaped profile. The only problem is I can't explain

how I get the spicules.

Thomas: Plus explaining all the other things

I outlined a minute ago: optical depth effects, collisions, etc.

Underhill : Right. I've somehow got to control

my spicules. But what we do know, without any doubt,

is that there is a certain amount of low density material leaving the star as a spherically expanding

shell. That is best shown by C III X5696, which is

excited by particular processes as I described in

the last session. It has a weak absorption component. The width of the emission gives the expansion

269
velocity. Then you can consider a completely unre- lated He I line, whose absorption edge is formed in

that part of the envelope between us and the shell,

and the displacement of this edge gives you the same expansion velocity. To interpret the velocity field

in the material forming the rest of the spectrum,

you need the trained field of spicules.

Payne - Gap o s ohkin : There are two reasons why

the material that produces the Gaussian profiles

must be near to the star: the first is that none of

these lines are forbidden, and the second is the

evidence from eclipsing stars.

Sohmidt-Kaler : I have the impression that the

profiles are separated according to excitation:

flat- topped lines, low excitation; Gaussian lines,

high excitation. Is this correct?

Undevhill : The separation is more by process

of excitation than by level of excitation. The

flat-topped profiles appear in lines that are ex- cited by particular processes. They're the lines

you see in emission in Of stars. I mentioned that

C III X5696 is excited by He II emission; so you

need He II A 30 3 and a supply of C"*" ions, which

implies a cool shell. In other stars N IV X4058

is flat-topped. For this you need collisional ex- citation: He+ ions colliding with N ions. He I X5876 may be flat- topped and possibly A3888; X4471

is too blended to be sure. The He I spectrum is,

I think, a simple recombination spectrum in an outer

part of the expanding, low density atmosphere. I emphasize again that you see these flat tops in only

one or two lines that can be excited in a low density

gas, and that their widths all indicate the same ve- locity as that shown by the few absorption edges.

Schmidt-Kaler : This would seem to confirm Kuhi '

model of an inner turbulent region and an outer ex- panding region: Flat- topped profiles are formed in

the outer region and Gaussian profiles in the inner

region.

Smith: Then it follows that the expansion ve- locities are greatest in the outer parts of the en- velope, and that the narrow lines - those from the

higher stages of ionization - are formed in the

innermost regions. This would imply that the ve- locity increases and ionization decreases outward.

Sahade: In Part B , we noted the existence of

two envelopes, at least in binaries. One lies around

the WR star; it is thick, it is where most of the

emission lines are formed, and it follows the motion of the WR star in its orbit. All the evidence

suggests that the matter in such an envelope is

270
being decelerated outward, implying that lines like

those of He II are formed close to the surface of

the star, and lines corresponding to higher degrees

of ionization are formed further out. Finally we reach the large expanding envelope where the lines

show dilution effects, and the radial velocities do

not follow the orbital motion of the star but always

show the same velocity.

Thomas: I too prefer this picture to Lindsey's.

Note that we have an initial jump in Tg from the

photosphere to the region of line emission, so initially at least the ionization must increase outward. This follows the behavior in the outer solar

atmosphere where again we have a region - the chromosphere and lower corona - of high excitation which

shows no evidence of expansion. The solar wind ex- pansion occurs much farther out, and indeed in a re- cent speculative paper on x-ray emission from WR

stars, Wallerstein and collaborators {Ap . J. 151 , L121, 1968) suggest the possibility of a stellar

wind. So all this accords with Sahade's suggestions.

Smith: Is there no evidence of an outward

increase in velocity in the solar chromosphere-corona?

Is it possible that the temperature increases and the

velocity decreases to a certain point where the situation then reverses?

Thomas: There is no evidence of expansion in

the solar chromosphere nor, indeed, in most of the corona. I think that for WR stars we can only try

to extend the work of Wallerstein and his collaborators. But if the expansion arises from the same source as the solar wind, rather than from some internal ejection process such as is usually assumed

for WR stars, we must know the value of Tg in the

WR corona. Maybe we have simply been arguing about

the wrong kind of ejection-expansion all these years.

Steoher: I have a point that serves as a transition between questions (2) and (3) . For y Vel we have two values of g corresponding to two

values of radius, which depend upon the assumed distance. The range in log g is 1.5 to 2.5. Given a photospheric temperature of 4x10 '^"K, can you con- struct a stable model in hydrostatic equilibrium

with a value of g in this range?

Underhill: A very strong no. With Tg = 3xlO''°K

and log g = 4, you can barely hold a model in hydrostatic equilibrium, particularly if you include the

radiation pressure in the lines.

Payne-Gap o s chkin : But you can't make a stable

model even for Rigel.

Underhill : Not one in hydrostatic equilibrium.

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Steoher: It is losing mass. I see Dopplershifted resonance lines both in absorption and in

emission in all supergiants.

Payne-Gaposohkin : And there are very interesting changes in the velocities of the absorption com- ponents of Ha, which are superposed on the emission

lines. Do I understand correctly that you cannot

make models for any of these stars?

Underhill : There are some model atmospheres

published at effective temperatures below 2x10** °K

with log g as low as 1. These are numerical calculations that I don't think have anything to do with

supergiants

Thomas: We turn now to question (4): What is

the complete structure of a WR object, and how does

the structure produce the atmospheric features?

J. Cox: I have one comment and one question

which is essentially the question Paczynski asked

in his paper. If one computes a period for radial

pulsations in the fundamental mode for pure helium

stars of one solar mass, the period turns out to be

about 10 to 15 minutes. For pure heliiam stars of

10 solar masses, the period is between 1 and 2 hours.

Now if the helium star has a hydrogen-rich envelope

on top of the helium core, the period will be longer

by some factor which will depend on the radius of

the star. The point is that the period could be

anywhere from a few minutes to perhaps a day. Now

Paczynski 's question is: Are light variations with

periods ranging from a few minutes to a few hours

observed in the WR stars?

Smith: Many stars are known to be erratically

variable. If the periods are that small, the periodicity could easily be missed.

Kuhi : 1 think the real answer is that no one has looked for such short term variations, although

it would be easy to do so. Merle Walker has detected a period of 71 sec in DQ Here, and I have recently been looking for a period of 1.337 sec in pulsars.

A moderate-size telescope of 50 or 60 inches could

easily do it, but we would have to be very careful

to distinguish between emission lines and continuum.

Presumably the variations we expect are in the continuum, so we would have to go to narrow-band work,

which means increasing the size of the telescope,

but it is still a soluble problem.

Castor: Do you know from the observations how

large a variation could exist?

Kuhi: HD50896, for example, varies 10 percent

or so over about a day. I don't know what causes

this particular variation. We could try to fit it

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as a binary, but it is not clear what is happening.

Nariai: I would like to point out that whatever the mechanism, the period must lie between 10^

and 10** sec: 10** sec is the diameter of the envelope

divided by the velocity; 10^ sec is the energy con- tained in the envelope divided by the energy flux

in the lines. So suppose we do find this period;

we cannot then conclude that the variation is really

due to the nuclear instability.

Thomas: You have apparently taken 1000 km/sec

as the velocity, and although this is a bit high even

for 1q 10 ^°K, that will only increase your upper

limit on the period. What would you buy for these

periods, John?

e7. Cox: The period depends on the size of the

envelope. 10** sec would be a reasonable order-ofmagnitude estimate for a large envelope.

Castor: If the envelope were so large that the

pulsation could be a progressive wave in the outer

part, the period should be dominated by the inner

part in which the pulsation is a standing wave. So

the period shouldn't be very long.

Hjellming: Wouldn't one expect two different

classes of optical evidence for pulsation? On the one hand, if the pulsation were spherically symmetric,

we might expect some periodicity in luminosity. But

if there are other kinds of pulsations, such as a non-radial pulsation, we might observe only statistical fluctuations in luminosity. I would think it

would be the latter alternative that would give us most of the evidence about the real situation. We

might also have harmonics of the radial pulsation.

J. Cox: I certainly agree that it might be

hard to observe non-radial pulsations directly; the

scale of the motions would be considerably smaller

than the size of the star. But if the pulsations

are generated by nuclear sources in the interior,

I think it unlikely there would be harmonics of the

radial pulsation.

Sohmidt-Kater : I believe Ross has made broad- band observations of about a half-dozen WR stars

and has found no variation. So I think this underlines Kuhi ' s remark that narrow-band work must be

used in the search for evidence for pulsations.

Underhill : For my 1968 Annual Reviews article

I collected what evidence there was from broad-band

work. The results are contradictory: One group will

say a given star varied, others will say it didn't.

Lindsey Smith has some evidence for variability from

her narrow band observations.

Smith: Not much evidence: Most stars reproduce

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from one night to another to within 0.01 to 0.02 mag.

So if variations are always there , they must be pretty

small.

Wrubel: If thermal instability is responsible

for some aspects of the WR phenomenon, the energy

released during the thermal pulse may not go into

the pulsational modes. Thermal instability produces

a convective region immediately above the energyproducing shell. Perhaps energy is pumped directly

into the turbulent zone. I would also remind you

that if, as Lindsey suggests, you want to add an additional parameter called "binary character" to

the physical properties of the star, you must really

add several parameters. The results on mass exchange

depend on the stage of evolution at which the mass exchange takes place. Thus several parameters such

as mass and separation of the components are involved.

Thomas: Are there any more comments on the

mass-exchange models? What about the problem of

chemical -differentiation, for example?

Wrubel: That aspect is the most speculative

part.

Thomas: Well, if you don't want the differences

between the WC and WN sequences to arise from ex- citation effects, you have a choice: Either you buy

initial differences in composition or you buy some- thing in the evolutionary process that produces these

differences. And Paczynski's remarks on chemical

differentiation are no more speculative than his argxjments on the production of supernovae. Actually

it is a bit harsh to refer to these suggestions as speculation. In a sense they are an extrapolation

to the next stage of the models from that sequence

which carries the evolving object up to the stage

of mass ejection with He- or C-burning cores and the

resulting instabilities.

Smith: It is on this basis that I prefer the

alternative I mentioned of C-enrichment of the atmosphere. On the basis of the Kippenhahn-Weigert series

of models we know we will get He-burning or C-burning models with a thin H envelope. We know they will

be unstable, and we know they will develop thermal

pulses. Thus it is likely we will get C-mixing in

the atmosphere. Agreed we are arguing on the basis

of extrapolation; I just think the distance we have

extrapolated is a bit less here.

Schmidt-Kaler : I agree that these mass-exchange

mechanisms seem quite well established. However, I think there may be two instability mechanisms working

to produce the WR phenomenon. First, the WR binaries

we see belong to the narrow-line, strong continuum

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class (Hiltner-Schild , WN-A) . Thus you have not an enhancement of the WR characteristics but the re- verse; the single stars seem to show the broader

lines, hence enhancement of the WR characteristic.

Second, all the WR stars that show the ring-like

shells seem to be single with the possible exception

of the weird star HD50896. A third rather weak piece

of evidence is that the Of stars do not seem very far

evolved. Fourth, the supergiants showing the same kind of high ejection velocities in the rocket UV are all single stars. So we seem to have strong evidence

for the existence of a second mechanism which doesn't

depend upon mass exchange.

Smith: In my summary, I said that another mechanism was probably needed for the single stars; and

the statistics do indicate that the binaries are younger than the single stars. Sco OBI is a good ex- ample, so I think this is virtually established.

Undevhill : I personally would like to see all

these suggestions on instabilities more firmly es- tablished before deciding on the evolutionary history of any of these stars. We know we require a source of mechanical energy and that this in turn

requires an instability. Several such instabilities

have been proposed, but only proposed. We do not

have a single definite calculation of how much energy any one of them actually contributes to the atmosphere . Thomas: Yes, Anne, but from the types of instabilities suggested and from the kinds of velocity

fields they are likely to produce, we now have something definite to calculate. The situation has improved enormously since 1938 when there was only a vague feeling that ejected shells had something to

do with WR phenomena; there was then no thought of

modifying the classical, radiation-dominated atmosphere. The situation has improved further since

194 8, when the idea that we needed an atmosphere

strongly influenced by mechanical supplies of energy

and momentum was only an assertion, and we had no idea where to get it. Also since then, the arguments have been stimulated by the Paczynski-Kippenhahn alternative to differences in excitation or initial composition as the explanation of the WC-WN

dichotomy. Even those objects which changed from

WN to WC may be telling us something. All these are definite conceptual models that we can investigate;

they are not just hypotheses about radiation pressure, magnetic fields, super-flares and super-spicules

Beats: If there were any truth in flare models,

you could not say there would be any one temperature

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at a given level in a WR star. There might be local

patches where Te is 10 times higher than in the sur- rounding areas. This could produce marked differences in the UV spectra of WR and 0-stars.

Thomas: Beals has raised the question of the

effect of inhomogeneities in the atmosphere. And

again we might turn to the Sun as a guide. People

are trying to explore the effects of inhomogeneities

in the visual by studying observations such as those

on granulation. They are also trying to differentiate between plage and non-plage regions by looking

at the rich variety of rocket-UV lines. So maybe

the results for the Sun will give us some insight

into Beals' suggestion.

Pecker-Wimel : Wray has shown me satellite

photographs of many supergiants and WR stars, so we are not limited to the single observation of y Vel that West says is all he can get from a rocket.

Wray: All we have so far are very low dispersion observations. We did observe several WR stars:

HD156385 and HD192163. The Apollo S-019 experiment

should give us the kind of observations discussed

here. There is a magnitude limit of about +9.5 m^

at A1350 with a dispersion of 60 A/mm.

Westerlund : I have one question with regard

to Paczynski's theory. In his paper V he says that

the companion should be a normal main-sequence star.

I believe that Kippenhahn and Lindsey expect the

same. But in the Galaxy there are at least 4 or 5 components classified as supergiants or evolved stars.

There are three in the Magellanic Clouds, and in

these three it is likely that the WR component is

WN7; you can see how the luminosity increases with

the luminosity of the supergiant. This does not seem to me to agree very well with the binary theory.

Kuhi : One might say something about the spectrum of the companion. Conti and I have started a program to look at the 0-type companions, and so far we have not been able to find any differences whatsoever in chemical composition or in anything else,

compared to ordinary O and B stars.

Westerlund : So you are suggesting that the

classification is wrong in those cases. I have not

classified any myself, but in Lindsey 's catalogue

of galactic WR stars, there are several binary systems with one component of luminosity class I. If

the WR stage lasts only 4x10^ years, that is hardly

sufficient time for a normal star to accumulate mass,

evolve up and off the main sequence, and become a supergiant. But the simplest explanation may be that

the classification is wrong. I would only like to

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know whether those you observed were selected so that the companion was classified as la? There were

only 4 or 5 of these out of our 50 or more binaries.

The chances are that you have not yet observed any

of them.

Smith: Let me ask a question. The time spent

in the phases that we suggest are WR stars is 5x10^

years, maybe as much as 10^ if we include stages E

and F. How massive does the secondary have to be on the main sequence to evolve in that time? We can get a pretty massive secondary, so if it evolves

within 10^ years we can get a supergiant with the WR

star.

Undevhill : I think it would be well to look at

those classifications again. It is extremely difficult to classify stars when you have two spectra on top of one another. If you think the H-absorption

lines look a little bit thin and a little bit sharp,

you may say la. But they will obviously look a little

bit thin and a little bit sharp just because they

have the WR emission on top of them. So I think the

most doubtful thing is the classifications.

Kuhi : Right.

Wrubel: I think Kippenhahn assumed that the

material that left the WR star was immediately deposited on the other star. Actually the other star

is well inside its critical lobe, and the material

that comes off the pre-WR star can float around for

a considerable time, possibly producing a spectrum

quite contrary to what you would expect from an or- dinary single star.

Thomas: In addition to the spectra of normal

stars and of supergiant components of WR binaries

and of WR stars , we now have the complicated spectrum produced by material that doesn't know whether

it belongs to the WR star or to the supergiant.

We have a Visiting Fellow Program in JILA that

provides a sabbatical year for people who will come, we hope, to pose interesting scientific problems to

us. Wrubel will be a Visiting Fellow next year; he

has just posed such a problem. We invite those of

you who are interested to consider doing likewise.

Thanks very much for being with us this week.

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