N O T I C E 

THIS DOCUMENT HAS BEEN REPRODUCED FROM 

MICROFICHE. ALTHOUGH IT IS RECOGNIZED THAT 

CERTAIN PORTIONS ARE ILLEGIBLE, IT IS BEING RELEASED 

IN THE INTEREST OF MAKING AVAILABLE AS MUCH 

INFORMATION AS POSSIBLE 
•^ ^•_ ,^a __. ^22

-

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RF Project 762142/712834

Final Technical Report

(NASA -CLc - 17b356) CCNTINUEG UlIbAYICLEI 486- 152 19

S:UDIE5 61 SVME bE STABS Gr LAILL TYNE

Fival Tccht,lcal kel ort, 15 May lStO - 31

Auq. 1985 (Ohio State UDiv-, Cclumbus.) Uuclas

76 p HC AC5 /Mk AC1 CSCL 03A G^/69 U4847

CONTINUED ULTRAVIOLET STUDIES OF SOME Be STABS OF LATER TYPE

Arne Slettebak

Department of Astronomy

For the Period

May 15, 1980 - August 31, 1985

NATIONAL AERONAUTICS AND SPACE ADMINISTRATION

Goddard Space Flight Center

Greenbelt, Maryland 20771

Grant No. NAG 5-52 3

November 1985

The Ohio State University

Research Foundation

1314 Kinnear Road

Columbus, Ohio 43212

..

^4r

- i

1

io
from NASA CP-2238, "Advances in Ultraviolet Astronomy* Four Years of IUE

Research", ed. Y. Kondo, J. K. read, and R. D. Chapman, 1982, p. 579

ULTRAVIOLET SPECTRA OF SOME BRIGHT LATER-TYPE

Be STARS AND A-F SHELL STARS .01

Arne Slattebak

.Parkins Observatory

Ohio State and Ohio Wesleyan Universities

ABSTRACT

Anomalous ionization (C IV and Si IV) is seen in IUE spectra of Be stars

as late as B8, and occurs also in standard stare of similar spectral types.

Asymmetrical lines suggesting mass lose are present in all the Be stars and

several of the standard stars as well, with no obvious correlation with

v sin i. Emission shoulders are present in the Mg II lines of two B ye stars

but not in Be stars of later type. Again there is no correlation with

v sin i. The A-F shell stare show rich Fe II absorption spectra in the

ultraviolet, in one case with velocity structure.

I. INTRODUCTION

Early ultraviolet spectroscopic observations of Be stars (e.g.

Marlborough ind Snow 1976) concentrated on the stars of early type,

demonstrating mass-loss effects in their spectra. More recently, these

investigations have been extended to cooler Be stars. Thus, Snow (1981)

found mass-loss effects in the spectra of Be stars as late as B6, and

Marlborough and Peters (1982) detected anomalc us ionization (C IV and Si IV)

in Be stars as late as B8. This investigation concentrates on the

ultraviolet spectra of the later-type Be stars and A-F shell stars, with

emphasis on mass-loss effects and anomalous ionization.

II. THE PROGRAM

High-dispersion IUE spectra of 13 bright later-type Be stars and A-F

shell stars were obtained in both the short and long wavelength regions, pus

8 standard stars of corresponding spectral types. These are listed in Tables

la and 1b, with spectral types and rotational velocities from a recent study

of the brighter Be stars (Slattebak 1982). The program and standard stars

were selected to include both sharp-lined an8 broad-lined stars of various

spectral types between 33 and F0.

Special attention was paid to resonance lines, for possible mass-lose

effects, and lines which indicated anomalous ionization in earlier

investigations. These included Si II 1190.413, 1193.284, 1260.418, 1264.730;

Si III 1206.510; C II 1334."32, 1335.708; Si IV 1393.755, 1402.769; C IV

1548.202, 1550.774; Al II 1670.787; Al III 1854.716, 1867.790; Fe II

multiplets 1, 62, and 63; and Mg II 2795.523, 2802.698. Some of these are

illustrated in Figures 1-4.

579
III. RESULTS AND DISCUSSION

A. The Be Stars of Later Type.

The approach taken in this investigation is that of spectral

classification: standard non-emission-line stars with MK spectral types

close to those of the program stars were observed along with the program

stars for direct comparison purposes. Furthermore, both sharp-lined and

broad-lined standard stars were included, to match as closely as possible the

sharp-lined (interpreted as pole-on, according to the rotational model) Be

stars and broad-lined (absorption-line shell.-type) Be stars. An unexpected

result is that the standard star ultraviolet spectra are almost as

interesting as the Be spectra, from the point of view of mass loss and

anomalous ionization.

1. Anomalous Ionization

According to Marlborough and Peters (1982), the coolest spectral type at

which C IV or Si IV would be expected under radiative equilibrium conditions,

either in the photosphere or in a cool circumstellar envelope, is BO and B2,

respectively. They observed both ions in WE spectra of Be stars as late as

B8, however. My results, obtained with different stars, confirm their

conclusion: the B8e and B7e stars 8 Cyg B and n Tau have both C IV (see

Fig. 1) and Si IV lines in their spectra. On the other hand, C IV seems not

to be present in the spectrum of the B5e star 8 Psc nor in the B6.5e

star 0 And (see Fig. 1), though both are pole-on stars like n Tau. Si IV

(see Fig. 2) is present in all the B5-6e stars as well as the later-type

stars shown in Fig. 1.

Among the standard non-emission-line stars, anomalous ionization effects

also exist. Thus, C IV is seen weakly in the B8 standard t And (Fig. 1) but

is very weak or not present in the B5 standards P Aur and t 2 Aqr. Si IV, on

the hand, is visible in the spectra of all three standard stars.

2. Mass-Loss

Fig. 1 shows that the C IV doublet is asymmetrical and shifted toward

shorter wavelengths in both n Tau and $ Cyg B. The same is true for the

shell stars 48 Lib, V Per, and o And. Fig. 2 shows that the Si I V doublet is

present and asymmetrical for the B5-6e stars shown there. The lines are also

asymmetrical for the B6.5-8e stars shown in Fig. 1. The Al III resonance

doublet is also strongly asymmetrical in the spectrum of ^ Per. Thus,

mass-loss effects appear to be present in the spectra of both shell and

pole-on Be stars of spectral types B3 to B8.

Interestingly, the standard non-emission line stars also show

asymmetrical lines, suggestive of mass-loss effects. Thus, Fig. 2 shows the

Si IV lines in the sharp-lined B5 standard P Aur to be strongly

asymmetrical, as are the Al III resonance lines in the broad-lined 35

standard ^ 2 Aqr. In Fig. 3, the Mg II resonance lines are clearly

asymmetrical in the spectrum of P Aur.

580

h

_^.r.... r+sti10
3. Emission Lines

Emission wings in the Mg II and Fe II resonance lines were observed in

nearly all the Be stars of earlier type studied with IUE by Dachs (1980).

Fig. 3 shows the resonance doublet of Xg II in the spectrum of a number of Be

and standard stars. Note that the lines have emission shoulders in the

spectra of both S Psc and • ^ Per. Both are of spectral type B5, but 5 Psc

is a pole-cc, star and * Per a shell star. The only other emission lines

detected in this investigation are the short-wavelength shoulders of the Fe I

multiplet 1 lines in V Per, but this is not a certain identification.

B. The A-F Shell Stars

The four A-F shell stars in Table la are identified as shell stars in

the optical wavelength region because of the simultaneous presence of

rotationally-broadened lines, from the underlying star, and sharp absorption

lines from ground states or low-lying metastable levels, presumably from some

kind of shell. All four show rich Fe II absorption spectra in the

ultraviolet (especially from multiplets 1, 62, and 63) which are considerably

stronger than in the standard star of corresponding type. Fig. 4 shows a

portion of the spectrum of one of the A-type shell stars, S Pic, with the A5

standard a :as. Note that the Fe II lines are not only stronger and deeper

in the shell star, but asymmetrical, indicating velocity structure.

It is a pleasure to acknowledge the very kind hospitality shown to me

during a stay at the Institute for Astronomy of the University of Vienna,

where much of this work was done. I am grateful also to Ted Snow for very

helpful discussions, Ken Carpenter for much help with computer aspects of

this investigation, Bob Wing for permission to use his spectra of S Cyg B,

Charlie Wu and Al Holm for their spectra of n U Ma, and the IUE team at

Goddard for their help and cooperation.

REFERENCES

Dachs, J. 1980, Proc. of Second European IUE Conf. (ESA SP-157. April :980), y

p.139. Y

Marlborough, J. M., and Peters, G. J. 1982, IAU Symp. 98, "Be Stars",

ed. M. Jaschek and H. -G. Groth (Dordrecht -Reidel), p. 387.

Marlborough, J. M., and Snow, T. P. 1976, IAU Symp. 70, "Be and Shell Stars",

ed. A. Slettebak ( Dordrecht -Reidel), p. 179.

Slettebak, A. 1982, Ap. J. Suppl. (in press).

Snow, T. P., Jr. 1981, Ap. J. 251, 139.

581
14b- 4_

r

1

TABLE la. PROGRAM STARS TABLE lb. STANDARD STARS

v sin l i v sin Ii

Name Sp. Type (km s ) Name Sp. Type (km s )

48 Lib B3:IV:e 400 n U Ma B3 V 150

y Per B5 IIIe 280 P Aur B5 V 90

9 Psc B5 Ve 100 1ZAgr B5 V 280

o And B6 IIIe 260 a Leo B7 V 280

And B6.5 IIIe 80 t And B8 V 90

n Tau B7 IIIe 140 a Lyr AO V < 10

B Cyg B B8 Ve 250 9 Leo A2 V 15

28 Tau B8 (V:)e 320 8 Cas A5 III-IV 100

1 Del B8-9e 280

8 Pic A5 IV-shell 120

11 Vul A5 IV-shell 200

m Leo A7 IV-shell 230 ..l,Sol

14 Cum FO III-shell 200

0

^,. nt. naol

. W IC OSI \.[x.01

r n L

y

L

VI

0

Y. N.=. x.01 hr Ht. (P01

o y

rw .^• i•m • W NL. ono)

I

i

o^

y1 ^ y uv I tool

of

Fig. 1. The C IV resonance doublet Fig. 2. The Si IV resonance doublet

in the spectra of B6.5-8 stars. in the spectra of B5-6 stars.

582

IV

4W
rr,

P r (alai «II Im

•. um

RM

•

move 1^

1 Apt

0

r. «a. 1.«I

1

... «.. I t.t1

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Rr VVW

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1

•

rw 17.• t/.• n.• '^ mile

Fig. G. Fe II multiplet 62 and 63

lines in the spectra of an A5 standard and an A5 shell star.

Fig. 3. (left). The Mg II resonance

doublet in the spectra of B5-6 stars.

583
Tin ASTROPHYSICAL JOURNAL Sum naNr SER= 33:869-892,1983 Decetrher

0 11063. The AwmwAn Asu0000ucal SOMM All .ighu raerwd Pnnted in U.S.A.

ULTRAVIOLET SPECTROSCOPIC OBSERVATIONS OF SOME Be STARS

OF LATER TYPE AND A-F TYPE SHELL STARS

ARNE SLETTEBA81 AND KENNETH G. CARPENTER

Perkins Observatory, Ohio State and Ohio Wesleyan Universities

Received 1983 April 12; accepted 1983 July 6

ABSTRACT

High-dispersion IU' :- mctra of 18 later type Be and A-F type shell stars plus eight standard

non-emission-line stars were analyzed for anomalous ionization and mass -loss effects. We find that

superionization exists in Be stars of the latest spectral subtypes (B8-B9) and is strongest in the Be

shell stars, but does not seem to be present in the A-F type shell stars. We also observe

superionization in normal B-type stars as late as BS and possibly B7-B8, but not later. Asymmetrical

or violet-displaced reswAnce lines, suggesting mass loss, were observed in all of our Be program stars

but one., and also in a number of our standard stars, but not in the A-F type shell stars. Lower limits

to the mass-loss rates computed from Si iv lines range between 5.3 x 10 -12 and 3.5 x 10 -11 Mo yr-1,

with die Be shell stars showing the largest values. Terminal velocities measured from the asymmetrical or violet-displaced resonance-line profiles show no clear correlation with ionization potential.

Mass loss is correlated with luminosity and effective temperature in the sense that the B-type stars

show raass-loss effects while the A-type stars do not, but seems to be uncorrelated with rotation.

Emis sion was observed only in the Mg u resonance doublet for two B ye stars in our sample. All of

the A-F type shell stars show strong Fe n and Mg n absorption spectra relative to standard stars of

similar spectral type. We conclude from our results that ( 1) although Be stars are nearly indistinguishable from normal B stars in the ultraviolet, there are differences in degree; and (2) the hot

.arcumstellar envelope surrounding Be stars which gives rise to superionized imes and mass loss

appears not to he as nearly equatorially confined as the cool circumstellar envelope.

Subject headings: stars: Be — stars: circumstellar shells — stars: mass loss — ultraviolet: spectra

I. INTRODUCTION

Early ultraviolet spectroscopic observations of Be stars

concentrated on the stars of early type (see, e . g., papers

in 1A Symposium 70 (1976n. More recently, these

observations have been extended to cooler Be stars.

Thus, Snow (1981) found mass-loss effects in the ultraviolet spectra of Be stars as late as B6, while

Marlborough and Peters (1982) and Doazan (1982) detected anomalously high ionization (superionization) in

Be stars as late as B7-B8. This investigation concentrates on the Be stars of later type, with observations

also of the A- and F-type shell stars, which are believed

to represent an extension of the Be phenomenon to later

apes. The object is to ascertain the degree to which

mass loss and superionization effects persist to later

types in the Be and A-F type shell stars, to compare the

ultraviolet spectra of the program stars with those of

normal, nonemission standard stars of similar spectral

types, and to search for correlations of the effects found

Guest Invesugatot, International Ultraviolet Explorer Satellite.

with stellar parameters such as u sin i. luminosity, and

effective temperature.

11. OBSERVATIONS

Fligh-dispersion IUE spectra of 16 program stars and

seven standard non-emission-line stars were obtained in

both the short (-1200-1900 A) and long (-1900-3200

A) wavelength regions during several US 2 shifts in

1980-1982. Archival IUE spectra for two additional

program stars and one additional standard star were

obtained from the National Space Science Data Center

(NSSDC).

The 18 program stars are listed in Table 1A in order

of spectral type, with spectral types and rotational velocities (u sin t) from the listing by Slettebak (1982). Stars

with narrow lines (small c sin i) were included, as well as

stars with broad lines (large u sin i ). The former may be

thought of as " poie-on" stars, if Be stars are interpreted

according to Struve's (1931) rotational model, while the

latter are then viewed nearly equatorially. These b;oadlined stars also often show "shell spectra" ( sharp

absorption lines which arise from ground states or meta869

^4
C+wj

870 SLETTEBAK AND CARPENTER Vol. 53

TABLE 1A

Trm PRoostAM STABS

V sin i Date of

HD Name/HR Spectral Type (kms -t) SWP LWR Observation

142983 ....... 48 Lib 83:IV:e-shell 400 11212 9829 1981 Feb 1

22192 ....... y Per B5 Ille-sbell 280 10388 9071 1980 Oct 17

217891....... d Psc B5 Ve 100 10386 9069 1980 Oct 17

217675 ....... o And B6 IIle-sbell 260 15508 11983 1981 Nov 15

183656' ...... HR 7415 B6e-shell 300 14478 11064 1981 Jul 15

6811 ....... ♦ And 86.5 Me 80 10387 9070 1980 Oct 17

23630 ....... 0 Tau B7 Me 140 10378 9060 1980 Oct 16

23862 ....... 28 Tau B8 (V: *Shen 320 10377 9059 1980 Oct 16

1839141 ...... je Cyg B 88 Ve 250 5369 4609 1979 D :1y 26

195325 ....... 1 Del B8-9:e-shell 280: 15514 11991 1981 Nov 17

88195 ....... 17 Sex Al V-shell 150 18819 14847 1982 Dec 19

42111` ...... HR 2174 (br) Al IV-shell 180 18821 14849 1982 Dec 19

112028° ...... HR 4893 Al IV-shell 200 18822 14850 1982 Dec 19

38090 ....... 12 Lep A3 III-shell 180 18820 14848 1982 Dec 19

39060 ....... # Pic A5 IV-sbell 120 11216 9835 1981 Feb 1

192518 ....... 21 Vul A5 IV-shell 200 15511 11988 1981 Nov 17

98058 ....... ♦ Leo A7 IV-sbdl 230 15504 11980 1981 Nov 15

108283 ....... 14 Com FO III-sbell 200 15506 11982 1981 Nov 15

'HD 183656: ME spectra obtainedby Y. Kondo and J. Sahade.

°HD 183914: IUE spectra obtained by R F. Wing.

`HD 42111: TIs is the brighter component of a visual binary system (ADS 4749; 5.8: 6.9; 29" separation). The fainter component is a

normal AO V star with o sin i • 100 km s -t (Slettebak 1%3).

°HD 112028: This u the brighter component of a visual binary system (ADS 8682; 5.3: 5.8; 22" separation). The fainter component is

itself a bu:ary, showing a composite spectrum corresponding to two sharp-lined spectra of types AO V and A2 V (Slettebak 1%3).

TABLE 1B

THE STANDARD STARS

HD Name/HR Spectral Type

U sin i

(km s " t ) SWP LWR

Dau of

Observation

120315'...... A UMa B3 V 150 3170 2518 1978 Oct 1 and 27

34759....... p Aur B5 V 90 10389 9071 1980 Oct 17

21%88....... 'P2 Aqr B5 V 280 10385 9068 1980 Oct 17

87901....... a Leo 87 V 280 10379 9061 1980 Oct 16

222173....... t And B8 V 90 10376 9057 1980 Oct 16

172167 ....... a Lyr AO V <10 11213 9830 1981 Feb 1

97633....... 0 Leo A2 V 15 15505 11981 1981 Nov 15

8538....... d Cas A5 III-IV 100 11214 9831 1981 Feb 1

'HD 120315: 1UE spectra obtained by A. Holm and C.-C. Wu

stable levels and presumably are formed in an equatorial

shell)—examples at various spectral types were selected

and are listed in fable lA It should also be noted that

while the Be stars all show emission, at least at Ha, the

A-F type shell stars in Table 1A do not—they are

recognized as shell stars because of the simultaneous

presence of rotationally broadened lines, on the one

hand, and sharp lines arising from ground states and

metastable levels, on the other.

The presence of the We shell star 48 Lib in i list of

Be stars of latex types also deserves some explanation.

This star has been claisified as B5 in the literature, but

this classification was probably based on the attribution

of its rather strong Si a lines at 4128 and 4131 A to the

underlying, star rather than to the shell. The correct

spectral type appears to be close to B3, but we have

retained 48 Lib as a program star because of its very

interesting spectrum.

Table 1B lists the eight standard stars. These %ere

selected to match, approximately, the spectral types of

the program stars, and also to show a range of line

broadening from sharp to broad. The spectral types are

from Morgan, Abt, and Tapscott (1978), Morgan and

Keenan (1973), or Johnson and Morgan (1953). The
No. 4,1983 Be AND A-F SHELL STAR W OBSERVATIONS 871

rotational velocities (u sin r) are from Slettebak er al.

(1975) or are recent determinations by us based on that

system.

III. SUPERIONIZAIION

An important result of ultraviolet observations of the

hotter stars was the discovery in the spectra of some oand B-type stars of atoms in higher stages of ionization

than would be expected if normal photospheric radiation mere the only ionizing source. Thus, lines of 0 vi,

N v, C rv, and Si Iv were observed in the spectra of stars

with photospheres too cool to form these ions, calling

for an additional source of ionization. This effect, usually called "superionization," has been reviewed recently

by Marlborough (1982) and Doazan (1982).

a) Be Stars and A-Type Shell Stars

Examples of superionization in the Be stars have been

reported by a number of investigators, including Doazan

(1982), Henize er aL (1976), Lamers and Snow (1978),

Marlborough (1977), Marlborough and Snow (1976),

Morton (1976, 1979), and Panek and Savage (1976).

Recently, using WE spectra of 23 be Stars with spectral

types 09-88 and luminosity classes III-V, Marlborough and Peters (1982) reported the probable presence

of N v lines in Be stars as late as B3, C iv lines as late as

B8 (B7, according to Slettebak 1982), and Si r y lines also

as late as B8 (also B7, according to Slettebak 1982).

They compare these observational results with theoretical estimates of the coolest spectral types (based on the

work of Lamas and Snow 1978) at which each ion

would be expected under radiative equilibrium con iigUrr (si RI slt 11501

. i

I i

FILAma

1

0

N W ss•a[.4 14W,

L

o '

Utap

^ssa ^4W '405 1410

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FIo. i

tions, either in the photosphere or in a cool circumstellar

envelope: 07 for N v, Bl for C rv, and B3 for Si iv.

One object of this investigation was to determine

whether superionization extends to spectral types later

than B7 in the Be stars and, indeed, whether the phenomenon reaches even into the A- type shell stars. Figtares 1 -6 illustrate the Si iv resonance doublet (multiplet

W 1) at 1393.755 and 1402.770 A, and the C iv

resonance doublet (multiplet W 1) at 1548.188 and

1550.702 A in the spectra of a number of the program

stars.

In Figures 1 and 2. both the Si iv and C iv lines are

strong in absorption in the spectrum of the B3 shell star

48 Lib. Note that there is some distortion of the profiles

by lines of Fe u from the shell. Figure 3 shows that the

Si ry doublet is clearly present in the three B5-B6e stars

shown there and is stronger in the two shell stars (> Per

and o And) than in P Psc. The behavior of the C IV lines

is consistent with that of Si iv: Figure 4 shows the lines

clearly in the spectra of the same two B5-B6e shell stars

but very weak or absent in # Psc.

Figures 5 and 6 show the Si iv and C iv doublets in

some B6. 5-B8e stars. The Si ry X1393.755 line seems to

be present in the three Be stars shown there, although

other lines distort the spectrum somewhat. The C IV

doublet also appears to be present in the spectra of n

Tau and P Cyg B, but is very weak or absent in the

B6.5e star,0 And. The presence of both Si iv and C iv in

the B8e star 0 Cyg B confirms the findings of

Marlborough and Peters (1982) and Doazan (1982) that

these lines persist into the late Be subtypes.

It is interesting that neither Si tv nor C iv is visible in

the spectrum of the B3 shell star 28 Tau (Pleione). This

p Wme ;c Z) 857 '150)

a LUX

rLUx

0

44 U► as. a . 1400)

1 '0^"N , -,

q n

0

IS45 550 555

,.zl

FIG. 2

Fto. 1.--Spectrai region including the Si tv resonance doublet at 1393 . 755 and 1402 .770 A in the 83 shell star 48 Lib and the B3 V

standard start UMa hest wavelengths of the Si Iv lines are indicated by arrows. the Fe It lines which distort the Si Iv profiles in the

spectrum of 48 Lib are at 1392 . 15, 139282, 1393 .21, 1401 . 77, and 1403.25 A.

Fto. 2.—Spectrid region including the C ry resonance doublet at 1548.188 and 1550.762 A in the W shell star 48 Lib and the B3 V

standard star q UMa. Rest wavelengths of the C ry lines are indicated by arrows. An Fe It line at 1550.27 A distorts one of the C Iv profiles.
p Au ► 131 )A1 131 (901

RELATIVE

FLUI( ?,,^ /U^

0

A,

►0e 131• (100)

L

p J J.

r.► 1520 (210)

,ky, i L

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0 ^

y' AA► 15 Y ( 2101

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1390 1395 1400 1403 1410

1,Ip1

FIG. 3.—Spectral region including the Si iv resonance doublet at 1393.755 and 1402 .770 A (rest wavelengths indicated by arrows) to three

85-136e stars and two B5 V standard stars.

^ a

e
Be AND A-F SHELL STAR UV OBSERVATIONS 873

P awr IC Is) 952 i9o)

RCLATIV(

FLUX

1

1

6570 (100)

0

Far 05 X 0 (200)

i 1

0

a And es a. ( 2110)

i 1

I

A 85X (2110)

1W 1550 '555

^1 a)

FIG. 4.—Spectral rcyop including the C tv rnonaace doublet

at 1548. 188 and 1550.762 A (rest wavelengths indicated by arrows)

in three B5-B6e stun and two B5 V standard stars.

object still has an extremtcly opaque shell, both in the

visible and in the ultraviolet regions of the spectrum (see

Figs. 16 and 21), which may affect the visibility of the

Si iv and C iv lines. On the other hand, the late-type

shell star 1 Del clearly ,bows Si iv A1393.755 absorption. Its shell is no, as dense as that of 28 Tau, but is

sufficiently opaqu e in the visible portion of the spectrum

to make the underlying star very difficult to classify.

The estimated spectral type of 1 Del is B8-B9, making

it one of the 1). test Be stars to show superionization

effects.

Marlborough and Peters (1982) have noted some dependence of the strength of superionized lines in their

B2e-B3e stars on v sin i. Visual inspection of their spectra revealed that the range of strengths of C iv lines

depends on v sin is while stars of large v sin i may have a

range of C Iv line strengths, stars of low v sin i seem to

have only weak C Iv lines. Although our sample is small,

our spectra would seem to confirm this result. Inspection of Figures 3-6 shows that Be stars with small v sin i

generally have weaker lines of Si tv and C iv than Be

stars of the same spectral type with large v sin i. A

possible interpretation, as Marlborough and Peters

(1982) point out, is that the hot component of the

circumstellar envelope in which the superiottized lines

arise is not spherically symmetrical.

Is superionization also a feature of A-type shell spectra? The answer seems to be no. Their is no trace of the

C ry resonance doublet in the spectra of our A-F type

shell stars. The Si iv lines are also probably not present,

although there is a dip in the -type shell star spectra

near the Si iv \1393.755 line. Figure 7 shows this

wavelength region for the Al iv shell star HR 4893, with

the AO V standard ac Lyr also shown for comparison.

The Fe n lines at 1392.15. 1392.83, and 1393.21 A are

prominent, and probably account for the observed depression. This conclusion is strengthened by the fact

that the other component of the Si iv doublet at 1402.769

A appears not to be present. Unfortunately, our spectra

are weak at these wavelengths, which makes the reality

of broad features difficult to ascertain.

Of course, superionization may occur in the form of

lines other than those of Si iv and C iv, and, indeed, in

the A-type shell stars we might expect less ionized

species than in the Be stars. In particular. if the Auger

ionization nimbaniam proposed by Cassinelli and Olson

(1979) is correct, the dominant stage of ionization in the

stellar wind expected under radiative equilibrium conditions is doubly ionized by X-rays. Unfortunately, the

resonance lino of C in, Si III, N ii-iv, and O n-v are

either outside the spectral range of 1UE or too weak to

detect in our .rUE spectra, but we were able to examine

the resonance lines of a number of other species:

L Si it multiplet UV 2 (1533.445 and 1526.719 A).

These lines increase in strength through the B-type stars
A),

874 SLETTEBAK AND CARPENTER

And lat aI aa4 M)

i 1

I

RELATIVE

FLUX

G

AM aa.S=a 160)

Vol. 53

ip Tau 973a (140)

I

0

a Crl a aaAa (2501

1 390 1395 1400 1409 1 410

X(i)

Fin. 5 --Spectral region including the Si tv resonance doublet at 1393.755 and 1402.770 A Irest wavelengths Indicated by arrows) in thr-e

B6 5-58e stars and a 118 V standard star.

and are very strong at A2, with w significant difference

between A-type shell stars and A-type standard stars.

Other WE spectra are too weak to show the lines for

types later than A2.

2. C n multiplet W 1 (1335.708 and 1334.532 A).

Our fUE spectra show these lines to strengthen from the

middle to the late B types and remain strong through A2

(again, our sensitivity is too low for detection later than

A2), with no appreciable difference between shell-type

and normal stars.

3. Al n multip.et UV 2 (1670.787 A). This line is

weak at B3, strengthens through the B types, and becomes very strong in the range B8-A5 (A5 being the

limit of our detection). The line is shown in AO-'k5 stars

in Figures 8 and 9, where again no significant difference

between shell-stars and standard stars is apparent.

4. Al to multiplet UV 1 (1854.716 and 1862.790 A).

These lines arc strong in the middle B-type staw as

shown in Figure 10. which also shows them . I , x

somewhat stronger and asymtretncal in the two Zell

stars ( q, Per and o And), indicating contributions from a

stellar wind. This will be discussed in the next section.

The Al to lines exhibit a general weakening from B8

into the late A types, with no striking differences be-
No. 4,1983 Be AND A—F SHELL STAR UV OBSERVATIONS 875

c and (C a) 99Y 1901 a Le A0 If (-,a)

IttLATrvt

FLUX

+ + i

nC.arrva

Rix

i

r^a_

,I a

0

And 84.5 tan (90)

of

 t^^q

1 1

1

It T9r 9TR. (140)

0

p cn a ag. Izso)

0

1545 1550 is»

a(1)

FIG. 6.—Spectral regsop including the C ty resonance doublet

at 1548.188 and 155C 762 A (rest wavelengths indicated by arrows)

us B6 . 5-Bbe stave, and a B8 V standard star.

tween the A shell and the A-type standard stars. The

lone exception is 0 Pic (A5 tv shell), which has rather

strong Al art lines relative to standard stars.

In summary, the various resatiz a lines we examined

in our WE spectra show no general tendency toward

supenonizauon in the A-type shell stars.

MR 4891 Ai x-OWN (too)

1790 799

!,(:1

FIG. 7. — Spectral region near the Si ry resonance line at

1393 . 755 A in an Al shell star and the AO V standard a Lvr (rest

wavelength indicated by an arrowl. The dip near the Si iv line

appe in to be due to the Fe a Una at 1392.15. 1392.82, and

1393.21 A.

b) The Standard Stars

An interesting result of this investigation is that some.

but not all. of the s=dard, non—emission-line stars also

show supenonizatioa effects in their ultra-violet spectra.

Figures 1 and 2 show that the B3 V standard i? UMa has

Si iv but not C iv absorption, as might be expected at

this specL"al type. However, both the sharp- lined ( p

Aur) and broad-lined ( ^, Aqr) B5 V standards show

strong Si iv absorption ( Fig. 3), while the C IV ,loublet is

weak or absent in both (Fig. 4).

For later spectral types it is very difficult to be certain

whether superonized species CWL The broad-lined B7

V standard a Leo appears to show a broad and shallow

Si iv X1393 . 755 line, but the other component of the

doublet is difficult to see. Figure 5 shows the depression

at the position of the Si iv X1393 . 755 line in the spectrum of . And. but it is distorted by the presence of a

nuw= of Fe a lines. Si iv X1402 . 769 seems not to be

presenL The C iv resonance doublet is not detectable in

the spectrum of a Leo but may be present in t And ( see

Fig. 6). Moving to still later types. Figure ' shows a

depression near Si ry X1393 . 755 in the spectrum of .ie

AO V standard a Lyr, but this is probably due entirely

to the confluence of the several Fe a lines in that

wavelength region. There is no trace of the C iv resonance doublet, nor does the A2 V standard B Leo show

either Si iv of C iv absorpuon.

0

0

ti

1
0

MR 2174 (Or) AI Z -shell (160)

o 

MR 4663 Al =-shell (200)

a Lyr AO !'10)

cI Ail

(-^ITtI

1

n

RELATIVE

FLUX

0

17 Sex Al 7- shell (150)

r-T'7-11

K .

r

Y

0

B Leo A27 (15)

I-1rT^1 +

1

I

1655 1660 1663 1670 1675 1680

X(A)

FIG. 8,—WaveImgtb region including the C I resonance multiplet (UV centered at 1657 A and the Al a resonance line (UV 2) at

1670.787 A in three Al shell stars and an AO V and A2 V standard star.
Be AND A-F SHELL STAR W OBSERVATIONS 877

12 Lop A3 3II -Shell 0801

CI AIA

rTl-Tt

I

RELATIVE

FLUX

0 L

q1

E Cal

^-lrTfi .

I

rrUN/

0

Pie

rlt-t-n +

0

1655 660 1665 1670

A5M - 1Z (100)

A51Z - Shill (120)

1675 1680

1,(A)

FIG. 9.,—Wavelength region including the C I resonance multiplet (UV 2) centered at 1657 A and the Al q resonance line (UV 2) at

1670.787 A in two A3—A5 shell stars and an A5 III—IV standard star.

We have already discussed resonance lines of other

species in the last section and found no strong evidenc:.

for unusual line strengths in either the program or

standard stars of later types. Our observations thus

show that superionization exists in standard, non-emission-line stars as late as B5 and, possibly, as late as

B7-B8, but probably not later.

IV. STELLAR WINDS AND MASS LOSS

There is now a considerable literature on stellar winds

and mass loss from Be stars. Marlborough (1982) and

Doman (1982) have summarized the ultraviolet data

which have been used to infer or demonstrate mass loss

in Be stars. Much of the work has centered on Be stars

earlier than B2 (e.g., y Cas, 59 Cyg, (' Tau, 0 Per). We

note here two papers which include also Be stars of

somewhat later type: Dacts's (1980) 1UE observation of

eight bright Be stars, and the work of Snow (1981) based

on Copernicus scans. The latter determined mass-loss

rates for 19 Be stars from analysis of Si In and Si IV

resonance lines, including objects as late as spectral tvpe

as B6.

Asymmetrical or violet-displaced resonance absorption lines, suggesting mass loss, were observed in all

(except 28 Tau) of our Be program stars and also in a

number of our B-type standard stars, but not in the

A-F type shell stars. This is shown for the Si IV and

C ry doublets in Figures 1-6. In addition to these lines,

stellar winds are also suggested in the profiles of resonance lines of Si Ii (multiplet W 5), Al. in (multiplet

UV 1), and Mg II (multiplet W 1) in some of our stars.

Figure 10 shows the Al III resonance doublet ^t 1854.716

and 1862.790 A in the spectra of three B5-B6e stars and

two Bf standard stars. Note the asymmetrical structure
where

d 1

w

T_ ire' f x0

me v„

878 SLETTFBAK AND CARPENTER

p Aur (A) Z) 857 (90)

0

Pso 8570 (100)

Pw 8 SIR 0 (M)

x

J

J

r

W

M

o

o And 56M0 (260)

in the shell stars t Per and o And, as well as in the

broad-lined standard 4, 2 Aqr. On the other hand, the

relatively sharp-lined (pole-on) Be star n Tau also shows

a slight asymmetry in the same sense.

The standard B5 V star p Aur also shows asymmetrical resonance lines of other ions, including Si 11 and

Mg a. The Mg a resonance doublet at 2795.528 and

2802.704 k in the spectra of several Be and standard

stars is shown in Figure 11, when the lines are clearly

seen to be asymmetrical in the spectrum of p Aur.

Attempts to determine quantitatively the rates of

mass loss associated with the winds from Be stars from

ultraviolet line profiles have been summarized by Snow

(1982x). A number of investigators, including Snow

(1981), have utilized the atlas of theoretical P Cygni and

stellar wind profiles by Castor and Lamers (1979). The

latter assume resonance scattering and use the Sobolev

approximation. They characterize the expanding en-

,.--lope by two functions: the velocity law v (r) and the

optical depth r(u). Thus, if v„ is the terminal velocity

and R. the photospheric radius, they define

w — v/v., x — r/R„

and choose a velocity law of the type

w- 0.01+0.99(1-1 /x)'8

(ft >01.

For an optical depth law Castor and Lamers adopt, for

the radial optical depth rim (w ), a law giving zero residual intensity at the violet edge of the line profile (i.e., at

W — — 1):

rim(„')— T(Y +1)(1—wo)- t- '(1—w)' (Y?0),

,y' an.. 857 (280)

Y

r

1850 1855 1860 1865

X(A)

FIG. 10.—Spectral region including the Al tit doublet at 1854.716 and 1862.790 A (rest wavelengths indicated by

arrows) in three B5-B6e stars and two B5 V standard stars.

f is the absorption oscillator strength, X 0 is the rest

wavelength of the transition, and N, the column density

of the absorbing ion in the envelope. Castor and Lamers

take wo — 0.01 for the flow speed at the base of the

envelope, choose # — 0.5, 1, 2, and 4, and point out that

the absorption part of their P Cygni profiles is very

insen&dve to the velocity law. Values of y range between C.5 and 4. They then give theoretical stellar wind

profile! for a variety of combinations of S, y, and T.

Sinc : the lines being studied may be formed in the

photosphere as well as in the winds of the stars under

consideration, it is necessary to correct for the presence

of underlying photospheric profiles. Castor and Lamers

suggest a procedure for this correction, which requires

that the photospheric profile be known. We follow their
993. (290)

I

AA

Aq ast (290)

( rJ 

M

1 AJY N ..1 . --A

t

n MII. _ + ',]'VnM ,]MV0 1 '

2790 2799 2900 2905

X (A)

Fin. I1.—Spectral region including the Mg it resonance doublet at 2795.528 and 2802.704 A (rest wavelengths indicated by arrows) in

three BS-B6e stars and two BS V standard stars.
880 SLETTEBAK AND CARPENTER Vol. 53

TABLE 2

Trz Ambac DATA

A E 1p IP

Ion (H - 1 .0) (A) f (ev)

Si ry .... 4.5 x 10' 5 1393.755

f

1402.770

0.528 45.14

0.262

Si 11...... 4.5 x 10 - s 1190.416 0.251 . 16.35

C tv .... 4.2 x 10 -4 (1548.188 0.194 64.49

11550.762 0.097

Al nr ... 3.3 x 10 - °

f 1854.716 0.539 28.33

11862.790 0.268

Mg a ... 4.0 x 10' s 12795 .528 0.592 14.97

17802.704 0.295

method here and adopt also the procedure by Snow

(1981), who derived the photospheric profiles directly

from the observational data. He assumed that the

long-wavelength side of a photospheric line profile will

be virtually unaffected by the wind if it is optically thin,

and folded the long-wavelength side of his observed

profiles over to obtain an approximation to the underlying photospheric profile.

Still following Snow's (1981) application of the Castor

and Lamers (1979) theoretical profiles, we calculated

our mass-loss rates from

M-1.74x10 -18r(w) y.R• IXOgAE'

where r(w) is calculated for w — 0.5, v„ is in units of

km s" 1 , R, is in solar radii, A is in angstroms, g is the

ionization fraction for the observed species, A E is the

assumed abundance of the observed element with respect to hydrogen, and the mass-loss rate is computed in

solar masses per year.

Table 2 lists the atomic data for the lines analyzed in

this investigation. The element abundances A E are solar

values from the compilation by Ross and Aller (1976),

while the wavelengths and I values are from Morton

(1978).

There are many uncertainties in th,; computation of

mass-loss rates from line profiles, some of which we will

discuss later in this section. The greatest uncertainty is

probably the estimation of the ionization fraction g.

This is because we have no complete physical picture of

how superionization arises; collision and radiative

processes may both play a role (see the discussion and

references in Marlborough 1982 and Snow 1981). In the

absence of definitive ionization equilibrium models, we

will assume that all of a particular atomic species exists

in the form of the ion being observed; i.e., we will take

g — 1 in all cases. This means that our mass-loss rates

will be lower limits in every case, since the atom may

exist in other stages of ionization which could also

contribute to the mass flow.

Table 3 lists the computed lower limits to the mass-loss

rates for both the program stars and those standard

stars which show asymmetrical or violet-shifted resonance line profiles. The stars are listed approximately in

order of spectral type, from the hottest to the coolest.

The Castor and Lamers (1979) parametet # has the

value of 0.5 for all cases and is not listed in Table 3 with

ther other parameters y and T. The stellar rad i i were i

estimated from the papers by Code et al. (1976) and

Harris, Strand, and Worley ( 1963).

In addition to the unknown ionization balance in the

stellar winds, other uncertautties affect the computation

of mass-loss rates. These have been pointed out by Snow

(1981) and, in a different context, by Kurucz (1974).

One difficulty lies in the determination of the terminal

velocity v,,, which has considerable uncertainty because

the short-wavelength wing of the line profile merges very

gradually into the continuum. Snow (1981) points out,

however, that the error in estimating v„ is partially

cancelled out in the profile fitting process, with the

mass-loss rate not affected radically. Another possible

source of error is the blending of the short-wavelength

wing of Si iv ,11393.755 with lines of Fe u, as can be

seen in Figures 1 and 7. This effect would be especially

troublesome for the shell stars and may lead to values of

the terminal velocity which are too large and therefore

also lead to excessive mass-loss rates. The assumption of

spherical symmetry in the Castor and Lamers (1979)

analysis may also not be valid for winds from Be stars,

although the evidence is somewhat mixed at this time

(see references in Snow 1981). Ktrucz (1974), in a

sample spectral synthesis near the C tv resonance doublet for a B2 star, has pointed out that the continuum

level and equivalent widths will be affected by rotational

line broadening. Furthermore, he shows that the central

wavelengths of the rotated features are a function both

of the blending with many weaker lines and the rotational velocity. Wavelength shifts produced in this way

aright then be erroneously interpreted as Doppler shifts

due to mass motion in the stellar atmosphere.

Recognizi * hat our assumptions are the same as his,

and that the. _ is evidence (see Marlborough 1982;

Doazan 1982) that Be winds are variable with time, it is

nevertheless interesting to compare our results with those

of Snow (1981). For three stars in common (48 Lib, +

Per, and o And), our terminal velocities are within 25%

of his, and our mass-loss rates derived from the Si iv

X1393.755 line (adopting his ionization fractions) agree

with his to within a factor of 2. On the other hand,

Table 3 shows that our mass-loss rates as derived from

the C iv lines are systematically lower and those from

the Al in lines systematically higher than the Si iv

mass-loss rates for a given star. A partial explanation

may He in the unknown ionization equilibria for all three

elements. Also, we have already mentioned that possible

blending of Fe it lines with the short-wavelength wing of

^4I
i

No. 4, 1983 Be AND A-F SHELL STAR W OBSERVATIONS 881

TABLE 3

TERMINAL VELOCITIES AND LOWER LacTs To THE MASS-Loss RATES

T, v„ Lower Limit to M

Stu R./Ro M" (K) Line (km S "1 ) y T ofo yr-')

48 Lib ...... 6 -4.1 18,000 Si ry A1393.755 610 2 0.25 2.2 x 10 - "

l C iv 11548.188 l 490 1 0.25 5.2 x 10 -' =

Si ry 11393.755 530 2 0.5 3.5 x 10

Per ....... 6 -3.5 15,000 C iv 11548.188

1A1

450 1 0.5 8.5 x 10 - ' =

III 11854.716 420 2 0.5 2.2 x 10 -10

Psc ....... 4 -2.4 15,300 Si iv A1393.755 480 0.5 0.1 5.3 x 10 -'=

Si ry A1393.755 700 1 0.25 2.7 x 10 - `1 p Aur ....... 4 -24 15,300 Si n X1190.416 200 2 1 1.6 x 10-

Mg It X 2802.704 230 2 0.25 2.1 x 10 - =

2 4 -2.4 15,300( Si ry X1393.755 480 1 0.1 4.8 x 10 - 12

Al In,\1854.716 320 1 01 2.2x10 -'1

Sirv,\1393.755 740 2 0.25 2.7x10-'1

o And....... 5 -3.1 14,000 C ry A1548.188

1A1

410 0.5 0.25 3.1 x 10 - 'Z

In X1854.716 380 2 0.25 7.2 x 10 -"

HR 7415 ... 5 -2.5 14,000( Si rv,\1393.755 600 0.5 0.25 2.6 x 10 - 1)

`C ry A1548.188 300 0.5 0.1 6.8)- 10- 1J

And ...... 5 -2.9 13,700 Si ry X1393.755 570 0.5 0.1 9.4 x 10-

Si ry A1393.755 550 0.5 0.1 8.7 x 10 -12 <

9 Tau ....... 5 -2.7 13,200 C ry ,\1548.188 450 0.5 0.25 3.8 x 10-12

Al tti,\1862.790 230 1 0.1 2.8 x 10-" i

a Leo ....... 3.6 -1.6 13,400 Si ry ,\1393.755 530 0.5 0.1 5.8 x 10

0 Cyg B .... 3 -0.7 12,000/ Si iv X1393.755 590 0.5 0.1 6.0 x 10-12

l C ry X1548.188 260 1 0.5 1.4 x 10

i Del ....... 3 -1.0 11,500 Si ry X1393.755 600 0.5 0.25 1.5 x 10-"

Si ry ,11393 . 753 may have resulted in systematically

large Si iv mass-loss rates.

Based on computed mass-loss rates, Snow ( 1982b) L°° s ,LOW" LIMT)

found the three non-Be stars in his sample to be indistinguishable from the Be stars and suggested that the Be

phenomenon itself is not uniquely linked to the presence

y,

of steEw winds. Although our ms as -loss rates in Table 3 0 S

may be quite uncertain, there is some validity in com -

paring the stars listed among themselves, as Snow has

done. We then also find no striking differences between

the Be stars and the standard stars, with one exception: -110 =

the five Be shell stars as a group appear to have some- J

what larger mass-loss rate lower limits than the other

° SHELL STAR stars in the table. This is also shown in Figures 12, 13, 1

and 14. Obviously, more observations of both normal

stars and Be stars are necessary to see whether this -115

correlation is valid. -0 5 -10 -15 -20 -27 -50 - IS -40 • s

The terminal velocities v„ listed in Table 3 range M°°^

between 200 and 740 km s -t, somewhat lower than FIG. 12.-Lower limits to the mass-loss rates derived from urc

either the values measured by Dachs ( 1980) or these Si ry X1393 . 755 line for 12 Be and standard stars vs. their bolometMarlboro 1982 for hotter Be stars. Our nc absolute magnitudes. Dashed listed by ^ ( ) to the data

line represents a least-squares fit
SLETTEBAK AND CARPENTER

LOO A (LOWER LINNT)

e

-IO.S

-II 0 •

e SNILL STAN

• • N

• STAIpAIO

-11 S

20,000 0

T• I-N)

Vol. 53

• o • ,MILL STAG

• STANDARD

100 200 300 400

r AM I I MT/••c )

882

L00 a ILOW[N LIMIT)

.(Os

8

-n 0

-ILS ^- IO,000 15,000

. _r+j

Fto. 13. —Lower limits to the mawleas rates derived from the

Si tv A1393.755 line for 12 Be and standard stars vs. their effective

temperatures.

Fto. 14.—Lower limits to the mass-loss rates derived from the

Si tv A1393.755 line for 12 Be and standard stars vs. their rotational velocities (v sin i ).

measured terminal velocities show no clear correlation

with ionization potential. In general, v„ is largest for

Si rv, followed by C iv, Al III, Mg u, and Si u. Three Be

stars in particular (> Per, o And, and il Tau) show

simultaneously asymmetrical resonance lines of Si rv,

C rv, and Al III, with terminal velocities decreasing in

that order. This result is not consistent with the usual

assumption that ionization increases outward in which

case the C ry lines should show the largos; terminal

velocities. Again, however, the Si 1v terminal velocities

could be systematically high due to blending of Fe u

lines with the short-wavelength wing of Si r y X1393.755.

We may attempt to correlate our mass-loss rates with

various stellar parameters. Snow (1982b) has discussed

the extension of OB star winds to lower luminosities,

finding that luminosity is the dominant factor in driving

the mass loss from Of to mid-B main-sequence stars.

Specifically, he believes that the winds in Be stars may

represent a straightforward extension of the OB star

wind phenomenon to lower luminosities. Insofar as we

find stellar winds in the Be stars of later type but not in

the A—F shell stars, which are of lower luminosity than

the Be stars, our results are consistent with Snow's

suggestion

Is there a correlation between the mass-loss rates and

the luminosities of our B-type stars? Figure 12 shows

our lower limits to the mass-loss rates derived from the

Si ry A1393 . 755 line for 12 stars versus their bolometric

absolute magnitudes. The latter were estimated from the

calibrations by Blaauw (1963) and Keenan (1963) and

the bolometric corrections of Code et al. (1976) and

Flower (1977), and are listed in Table 3. Again, a

correlation seems to be present in the sense of Snow's

(1982b) conclusion: the more luminous stars show

somewhat larger mass-loss lower limit,,. The dashed line

in Figure 12 represents a :east-squares fit to the data

and has a correlation coefficient r — 0.50.

A correlation between mass loss and effective temperatures also exists in the sense that mass-loss effects are

found in the B-type stars but not in the A-type stars.

Within the B-type stars which show mass-loss effects,

however, no correlation with effective temperatures is

evident. This is shown in Figure 13, in which our

mass-loss lower limits derived from th, Si r y X1393.755

line are plotted versus effective temperatures for 12

stars. The latter were estimated from the papers by Code

et al. (1976) and Flower (1977), and are listed in Table 3.

A straight-line least-squares fit to the data in Figure 13

has a correlation coefficient r — 0.26, which we consider

not to be significant.

Figure 14 shows our lower limits to the mass-loss

rates derived from the Si iv X1393.755 line for the same

12 stars plotted against their rotational velocities v sin i.

Here we might expect a correlation for the Be stars if

stellar winds are related to the Be phenomenon, since

the latter is clearly related to stellar rotation. No correlation is evident: a straight-line least-squares fit to the

data in Figure 14 has a correlation coefficient r of only

0.26.

Plots of bolometric absolute magnitudes, effective

temperatures, and rotational velocities for the B-type

stars versus lower limits to their mass-loss rates derived

from C Iv and Al III lines were also made. Although

fewer observations were involved, these generally showed

the same trends as the Si ry plots: no correlations of

mass loss with effective temperatures and rotational

velocities, but a mild correlation with bolometric absolute magnitudes.
.A.

No. 4, 1983 Be AND A-F SHELL STAR UV OBSERVATIONS 883

V. EMSSION LINES

Be stars are characterized by Balmer emission in the

visible portion of the spectrum, but show little or no

emission in the ultraviolet. In a study of eight bright Be

stars, selected because their optical spectra show very

strong Balmer line emission, Dachs (1980) found emission is WE spectra only in the wings of the Mg u

resonance doublet and (for two stars only) in the wings

of the Fe u resonance multiplets UV 1 and UV 2.

Brubweiler, Morgan, and van der Hucht (1982) observed

Mg u k and k emission in the spectra of the early-type

Be stars # Per, v Cyg, and 7 Cas using the balloon-borne

ultraviolet stellar spectrometer (BUSS), and suggested

that a Bowen mechanism driven by Ly P is responsible.

Only two of our program stars show definite emission

in the wings of the Mg n resonance lines: P Psc and ¢

Per. These lines are shown in Figure 11. It is interesting

uxat both are B5 stars, but t Per is a shell star with a

large v sin i, whereas ^ Psc has a small v sin i and is

presumably a pole-on star. The Mg n resonance-line

region in the spectrum of the B3 shell star 48 Lib is

shown in Figure 15, along with the corresponding wavelength region in the standard B3 V star n UMa. Notice

that the Mg n lines in 48 Lib are deep and show velocity

structure, suggesting motions in the shell but no obvious

emission. Dachs (1980) also found Mg u emission in 48

Lib to be uncertain.

In order to check for possible variability of the Mg n

emission features in the spectra of the aforementioned

stars, additional IUE spectra of ^ Per (LWR 9832 on

1981 February 1 and LWR 11 990 on 1981 November

17) and 0 Psc (LWR 11989 on 1981 November 17) were

obtained. No significant differences are discernable in

these spectra, indicating no major changes in the Mg u

emission over time periods of the order of a few months

to a year for these stars.

UN. (loll as 111001

1 i

11LAIM

^

0

NW \7•1Z•• 1400)

0

rr.o :rw

ssoo ^Iljl

:.°°

Mo. 15.—Spectral region ingluding the Mg it doublet at 2795.528 and 2802.704 A (rest wavelengths indicated by

arrows) in the B3 shell star 48 Lib and the B3 V standard star q

UMa.

The Mg n resonance doublet in the spectra of other

program and standard stars is shown in Figures 16, 17.

and 18. There is no obvious emission for any of the

stars, and we conclude that Mg u emission does not

occur in spectral types later than B5 in our sample.

It is interesting that Blanco et al. (1982) find faint

chromospheric emission in the Mg u k line of a Aql, a

rapidly rotating A7 which is on or near the main sequence. On the other hand they find no Mg Ii emission

in HD 192518 - 21 Vul, which is one of our A-type shell

program stars. We also see no Mg a emission in the

spectrum of 21 Vul (see Fig. 18), nor is emission present

in any of our other A-F type shell stars.

Figures 19, 20, and 21 show the wavelength region

which includes the strongest lines of the Fe ti multiplet

UV 1 in some of our Be stars and standard stars.

Figures 22 and 23 show lines of the multiplets UV 62,

UV 63, and UV 64 of Fe u in the spectra of our A-F

type shell stars plus standard stars. The only object

which may have faint Fe n emission wings is the B5

shell star t Per (see Fig. 20, in which the short-wavelength wings of the Fe a multiplet UV 1 lines appear to

be somewhat higher than the long-wavelength wings),

but this is uncertain.

V1. THE A—F TYPE SHELL STARS

The eight A-F type shell stars listed in Table IA

range in spectral types from Al to F0. In the optical

spectrum, all show rotationally broadened lines (v sin i

ranging from 120 to 230 km s -1 ) plus sharp lines which

arise from ground states or metastable levels, the latter

presumably formed in some kind of shell around the

rapidly rotating underlying star. None of the stars show

Balmer emission, although the hotter ones may have

sharp Balmer absorption cores, which presumably also

form in the shell. While the hotter A-type shell stars

may also have sharp Fe ii lines in the optical spectrum

(arising from metastable levels 2.6 to 2.8 eV above the

ground state), the cooler ones are more likely to show

the Ti n lines at 3685, 3759, and 3761 A (arising from

metastable levels - 0.6 eV above :be ground state) and

the H and K resonance lines of Ca t:.

We have already noted that the A-F type shell stars

seem not to show superionized lines, nor do they have

asymmetrical or violet-displaced resonance lines suggesting mass loss. Figures 22 and 23 show the strongest lines

of Fe iI multiplets UV 62, UV 63, and UV 64 (with

lower metastable energy levels at -1 eV) in the spectra

of A-F type sheL' stars and standard stars. In Figures 8

and 9 we show the wavelength region around the Al u

rmnance line at 1670.787 A, including also the resonance multiplet of C i centered at 1657 A, for five

A-type shell stars plus three standard stars. We have

already noted that the A-F type shell stars seem not to

0•

i
i r

r_

r,

C+j

L And Img 11) 88 3r (90)

i

I

RELATIVE

FLUX 1.J

V

0

I - - r -VV`'Vvyvy

0

77 TOY 137 $ • (140)

1

+ +r

0 1

Q Cyg 8 B87s (250)

i r

i

28 Tau 88():)* (320)

I + A'1 i

0

2790 2795 2800 2805 2910

X (A)

FIG. 16.— Spectral region including the Mg it resonance doublet at 2795.528 and 2802.704 A (rest wavelengths indicated by arrows) in

four B6.5-88e stars and an B8 V standard star.

884

dal
lir

a Lyr (MOII> AO 1.10)

I

RELATIVE

FLUX

0

17 Sea Al 7-shell (ISO)

0

MR 2174 Isr) Al M-sMll (180)

0

MR 4893 Al lY -shell (2001

0 —+

B Loa A27 (15)

I

0

2790 2793 2800 2905 2810

1(41

Flo. 17.—Spectral region including the Mg if resonance doublet at 2795.328 and 2802.704 A (rest wavelengths Indicated by arrows) in

thm Al shell star and an AO V and A2 V suaadard star.

885

0
(Mg 111

I+

A5z-3z (100)

A512-Moll ( 120)

i

A517-shell (200)

+ 1

a Coe

1

RELATIVE

FLUX

0

wa

1

0

21 VYI

A7 X-ehell (230)

0

p Leo

1 i

0

k

. f+7

14 Cons F014 -eholl (200)

^N

0

2790 2795 2800 1 2805 2810

X!A)

Flo. 18.—Spectral region including the M6 It resonance doublet at 2795.528 and :802.744 A (rest wavelengths ; iatiatvd by 3rrowsl in

(our A5-FO shell star and an A5 III-IV standard star.

- J
Be AND A-F SHELL STAR UV OBSERVATIONS 887

iwa ^. fa 1 tN i usol

A&ArW

FLUX

1

u ^r sa a .aoo^

a

2!M rm 23M 20M

kill

Flo. 19.—Spectral region including the strongest lines of the Fe it multiplet UV 1 in the B3 shell star 48 Lib with the spectrum of the E13

V standard star q UMa shown for comparison. The Fe it at 2585.876 and 2399.395 A are indicated by : - ows.

show ultraviolet emission lines, as may be seen in the

above figures and, for the Mg n resonance lines, in

Figures 17 and 18.

As might be expected. the Fe Il lines which arise

either from ground states (e.g., multiplet UV 1) or from

metastable levels (see Figs. 22 and 23) are stronger,

deeper, and sharper in the A-F type shell stars than in

the standard stars. At least one star ( ft Pic) shows

structure in the Fe II line profiles, suggesting velocity

differences in the shell. Figures 17 and 18 also show that

the K. resonance lines ( analogously to the Ca It H

and k .es in the optical region of the spectrum) are

stronger and deeper in the A-F type shell stars than in

the standard stars.

VII. SUIAUSAIY AND CONCLUSIONS

We may SUMMAri7C Olit results as f0110w3:

1. Superionization in the Be stars extends to the latest

spectral subtypes (B8-B9) but does not seem to be

present in the A-F type shell stns.

2. The superionized lines in the Be stars appear to be

correlated with v sin i in the sense that they tend to be

stronger in stars with large v sin r than in stars with

small v sin t at a given spectral type. They are strongest

in the Be shell stars, which are also charactertred by

large rotational velocities. These results suggest that the

hot component of the circumstellar envelope from which

the superposed lines originate is not distributed with

spherical symmetry, as has been suggested earlier by

Marlborough and Peters ( 1983).

3. Supenonization is also observed in our sample in

normal, non-emission-line stars of spectral type as late

as BS and, possibly, as late as B7-B8, but not later.

There is no obvious correlation in the strength of the

supenonized lines with v sin i.

4. Asymmetrical or violet-dsplaced resonance lines, i

suggesting mass loss. were observed in all of our Be

program stars but one, and also in a number of our

standard stars, but not in the A-F type shell stars.

S. Computed lower limits to the mass-loss rates for

the Be stars are in the range 1.4x 10 10 - "-2.2x `0 M,-.

yr - '. The results from the Si IV ,1393.755 line range

between 5.3X10 - ' 2 and 3.5 x 10 - " M_ yr"'. Thee

numbers are consistent with results for Be stars in

common with earlier investigators.

6. The five Be shell stars as a group show the largest

lower limits to the mass-loss rates of the stars in our

sample.

7. The terminal velocities measured from the as ymmetrical or violet-displaced resonance-line profiles range

between 200 and 740 km s -', but show no clear correlation with ionization potential. This result is inconsistent

with the usual assumption that ionization increases outward.

8. Mass loss and luminosity are correlated for the

stars in our sample in the sense that the B-type stars

show mass loss effects while the (lower luminosity) Atype stars do not. Among the B-type stars themselves,

the more luminous objects statistically have larger lower

limits to the mass-1065 rates than the less luminous stars.

9. Mass loss and effective temperature are correlated

for the stars in our sample in the sense that the B-type

stars show mass-loss effects while the (cooler) A-type

stars do not. No correlation is found among the Btype stars themselves.

10. Mass loss and rotation appear to be uncorrelated

for the stars in our sample.
0

9 And 96 m 6 ( 260)

*` Aor 851 (260)

0

'i

p Aw 1F61) 99T tool

I

A[t,111VE

FLUX

0 l^

19 ►66 ` 9676 (100)

0

M 95x6 (2901

A A +

I

0

i

2565 2590 2596 2600

41 Al

FIG 20 —Spectra; region including thq strongest fines of the Fe it UV 1 m three B5-B6 • stars and two B5 V standard stars

The Fe it at :585 876 and 2599 395 A are indicated by arrows.

888
0 '

I

0

28 Tau

a c yl a more (250)

I -

11(Y )• (320)

(*-T

AM IF.II) GeV fool

P4LATIVIE

FLUX

0

AsId 11* (00.1

V!W

91 Tea 87 2 • 11401

A 11. Jkl Atli .111

0

25" 2590 25" ZGW

XIA)

FIG. .1 1.—Spectrall rgWon including the stroagat U= of the Fe it mulupict UV I in fOUf B6.5- BAt Stan and A EIS V %twdard Star The

Fe 13 Una at 2585.876 and 25" 395 A are ind-c-tcd by arrows.

889
o Lrr (FS II: AOY 1.10)

63 63 62 62 63 62.63 62 63 64

1 1 1 1 1 1 1 1 1

RELATIVE

vi y'vvy a

I,(^4411/ Ii/' J^

FLUX

0

17 Sol Al Y-seal (150)

! I I I I I I I I

1 ;

1NA^ 1 ^V1ry

' M h ^'y)

MR2`74( pr) AlM-sAal (180)

I ! I I I I I I I

i

1

i

U

MR 4893 Al M-sholl 12001

I I I I I I I 1 1

i

0

B Leo A27 (15)

I I i II I I II

'r

0

2735 2740 2745 2750 2755 2760 2'165

X(4)

FIG 22. — Spectral region including the strongest lines of the Fe a multiplets UV 62, UV 63, and UV 64 in three Al ,hell stars and an AO

V and A2 V standard star.

890

,+1

C

"v
0

Leo

i

I I

^ r .J Al l 1 u ^ . ^ ,

A7=-shall (230)

I i

0

14 con+

I i I ii i

FOM-shall (2001

S 

F

t

1

S Cat ( Fs II 1 AS A•$ U00)

63 63 62 62 63 62,63 62 63 64

I 1 { I I 1 I I I

RELATIVE

FLUX

v

Pic

•

'I

l 4 i I I i

1 r _k% . k J

AS M-shall (1201

I

1^

I 1

r l^y`ll,`

r 

0 `^

21 Vul

r I^I

I f I 1V i

61 L 1

AS S - shell (200)

I 1 I

0 J

2735 2740 2745 2750 2735 2760 2765

X(A)

FIG. 23.-- Spectral region including the strongest lines cf the Fe it multiplets UV 62, W 63, and UV 64 in four A5— FO shell stars and an

A5 III— IV standard star.

891

LIJW
a`

892 SLETTEBAK AND CARPENTER

11. The only definite emission we observe is in the

wings of the Mg n resonance doublet in two B5e stars in

our sample—no Be stars of later type nor any of the

A-F :ype shell stars show definite emission in any

ultraviolet line.

.2. All of our A-F type shell stars show strong

ultraviolet Fe u and Mg n absorption spectra relative to

standard stars of similar spectral type.

From these results, we draw the following conclusions:

i) Although Be stars are indeed nearly indistinguishat'e from normal B stars in the ultraviolet, as has been

pointed out by Snow (1982c) and Doazan (1982), there

are differences in degree: Be stars show superionization

to later specusl types, generally have stronger super

ionized lines at a given spectral type, and, at least for

the shell stars, show larger mass-loss rates.

ii) R correlations between the existence of superiopized lines and mass loss versus v sin i exist for Be

stars, they are weaker than correlations (from optical

spectra and polarization measurements) involving the

cool circumstellar envelope. This suggests that the latter

is more nearly equatorially confined than the hot envelope which gives rise to the supenonized lines and

mass loss.

We are grateful to Mike Marlborough, Ted Snow,

Paul Barker, and George Collins for very useful discussions. Some of this work was done while one of us

(A. S.) was a visitor at the Institute for Astronomy of

the University of Vienna, and the kind hospitality offered

there is gratefully acknowledged We thank also the IUE

team at Goddard Space Flight Center for their help and

cooperation, and the Ohio State University Instruction

and Research Computer Center for providing computer

time and fatalities. The many drawings in this paper

were prepared by Peter D. Stoycheff. We are also grateful to Mark Wagner for obtaining optical spectrograms

of 17 Sex for u& Finally, we thank NASA for support

via grant NAG 5-52.

REFERENCES

Blaauw, A- 1963, in Stars and Stellar Systems, Vol. 3, Basic

Astronomical Data, ed. K. Aa. Strand (Chicago: University of

Chicago Press), p. 383.

Blanw, C. Bruca, L., Catalano, S., and Marilli, E. 1982, Aim Ap.,

115, 280.

Bruhweiler, F. C., Morgan, T. H., and van der Hucht, K. A. 1982,

Ap. J., 262. 675.

Cassine& J. P., and Olson. G. L. 1979, Ap. J., 229, 304.

Castor, J. 1., and Lamers, H. J. G. L. M. 1979, Ap. J. Suppl., 39,

481.

Code. A. D., Davis, J., Bless, R. C., and Hanbury Brown, R. 1976,

Ap. J., 20, 417.

Daehs, J. 1980, in Proc. 2d European 1UE Conf. (ESA SP-157), p.

139.

Do— , V., 1982, in B Stan with and wit,iout Emission Lines, ed.

.4. Underhill and V. Do—n (NASA SP-456), p. 375.

Flower, P. J. 1977, .4srr. Ap , 54, 31.

Harris, D. L., III. Strand, K Ace, and Worley, C. E. 1963, in Stan

and Stellar S ystems, Vol. 3, Basic Astronomy -al Data, ed. K. Ace

Strand (Chicago: University of Chicago Pre :), p. 273.

Henze, K G., Wray, J. D., Parsons, S. B., aid Benedict, G. F.

1976, in IA U Symposium 70, Be and.'hell Stan, ed. A_ Slettebak

( Dordrecht: Reidel). p. 191.

Johnson. H. L., and Morgan, W. W. 1953, .4p. J., 117, 313.

Keenan, P. C. 1963, in Stars Stellar Systems. Vol. 3, Basic

Astronomical Data, ed. K Ace Strand (Chicago: University of

Chicago Press), p. 78.

Kunrea, R. L. 1974, App. J. ( Letters), 188. L21.

Lamers, H. J. G. L. M., and Snow. T. P. 1978. Ap. J., 219, 504.

Marlborough, J. M. 1977, Ap. J., 216, 446.

1982, in IA Symposium 98. Be Stars. ed. M. Jaschek and

H.-G. Groth (Dordrecht: Reidel), p. 361.

Marlborough J. M ,and Peters G. J. 1982 in IA U Symposium 98,

Be Stan, ed. Ni Jaschek and H.-G. Groth (Dordrecht: Reidel),

p. 387.

.Marlborough, J. Ni, and Snow, T. P., Jr. 1976, in 1A U Syymposium

70, Be and Shed Stan, ed. A. Slettebak (Dordrecht: Reidel), p.

179.

Morgan. W. W., Abt, H. A., and Tapscott, J. W. 1978, Revised MK

Spectral Atlas for Stars Earlier dim the Sun (Yerkes Observatory, University of Chicago and Kitt Peale National ObservaMorg n, W. W., and Keenan, P. C. 1973. Ant. Rev. Astr. Ap., 11,

29.

Morton, D. C. 1976, Ap. J., 20.3, 386.

1978, Ap. J., 222, 863.

1979, M.:Y. R. A. S., 189, 57.

Panek, R. J., and Savage, B. D. 1976, Ap. J., 206, 167.

Ross, J. E, and Aller, L. H. 1976, Science, 191, 1223.

Slettebak, A. 1%3, Ap. J., 138, 118.

1976, ed. IA U Symposium 70, Be and Shell Stan

( Dordrecht: Reidel).

. 1982, Ap. J. Suppl., 50, 55.

Slettebak. A.. Collins, G. W., II, Boyce, P. B.. White, N. M., and

Parkinson. T. D. 1975, Ap. J. Suppl., 29,137.

Snow, T. P. 1981, Ap. J., 251, 139.

Snow, T. P., Jr. 1982x, in IAU Sympasium 98, Be Stars, ed. M.

Jascbek and H.-G. Groth (Dordrecht: Reidel), p. 377.

1982J, Ap. J. ( Letters), 253, L39.

1982 c, in Advances in Ultraviolet Astronomy : Four Years of

1UE Research, ed. Y. Kondo, J. %i Mead, and R D. Chapman

(NASA CP-2238), p. 61.

Struve, O. 1931, Ap. J., 73, 94.

KENNETH G CARPENTER: JILA, University of Colorado, Boulder, CO 80309

ARNE SLETTEBAK: Perkins Observatory, PO Box 449, Delaware, OH 43015
THE ASTROPHYSICAL JOURNAL, 286:741-746,1984 November 15

1994 The Amencan Astronomical Smmy All nights rn:rvel Pnated to U.S.A.

ROTATIONAL VELOCITIES OF LATER B TYPE AND A TYPE STARS AS DETERMINED

FROM ULTRAVIOLET VERSUS VISUAL LINE PROFILES

KENNETH G. CARPENTER, ARNE SLETMBAK, t AND GEORGE SONNEBORN

Perkins Observatory, Ohio State and Ohio Wesleyan Universities

Received 1984 March 11: accepted 1984 June 6

ABSTRACT

Theoretical line profiles of the Si Rt 1299 and Fe it 2756 lines are computed for rotating B5-A7 model stars.

and compared with observed profiles from IUE spectra to derive rotational velocities. Real differences in

widths for ultraviolet as compared with visual line profiles exist in our sample of B type stars (but not for the

A type rapidly rotating stars), although these are not as large as previously reported in the literature. Comparison with our theoretical line profiles gives rotational velocities that are in good agreement with visually

determined v sin i values for the same stars, which suggests that our shape-distorted, gravity-darkened models

are reasonable.

:;ubtect headings: line profiles — stars: early-type — stars: rotation —• ultraviolet: spectra

1. IN , RODUCTION

Most recently, Marlborough (1982) reviewed and commented on the observed discrepancy between rotational velocities (v sin i) of Be stars as estimated from widths of absorption

lines in ultraviolet as compared with visual spectra. He pointed

out that perhaps the first to find such an effect were Morton et

al. (1972), who observed several ultraviolet absorption lines

(primarily C tt and C tit) in the rapidly rotating star Ophiuchi

to be unusually narrow. They suggested that the narrow lines

may be formed predominantly in the more slowly rotating

polar regions or that they may originate in a shell that has a

lower rotational velocity than the visual photosphere. Not long

after, Heap (1975, 1976, 1977) found ultraviolet lines of C In, C

Iv, Si in, and Si Iv in the spectrum of the shell star Tauri to be

relatively narrow, corresponding to a considerably lower rotational velocity than is estimated from visual line profiles.

Hutchings (1976a), in attempting to interpret Heap's (1975,

1976, 1977) observations, suggested also that in the farultraviolet spectrum of a star with a temperature gradient

across its surface, the underlying continuum radiation may

come preferentially from polar regions, where v sin i is expected

to be small. He followed this up (Hutchings i976b) with a

quantitative analysis, finding that v sin i from an assumed line

profile at 4500 A may be twice that at 1000-1200 A under

certain conditions. Sonneborn and Collins (1977) performed a

sirr_ilar but more refined analysis by assuming a rotationally

distorted star and including the temperature and gravity

dependence of actual lines. Their results were qualitatively

similar to those of Hutchings 11976a, b), but they found a

smaller variation in rotational line broadening than that predicted by Hutchings or observed in ,' Tauri by Heap.

There are problems in both the observational and theoretical areas, however, as Marlborough (1982) pointed out. In light

of more recent ultraviolet observations of a number of Be stars,

Heap's (1975, 1976, 1977) assumptions that all strong resonance lines in the ultraviolet are photospheric and that the

circumstellar envelope would not contribute to lines of high

stages of ionization are probably not correct. On the theoretical side, not only rotational shape distortion and the

Guest Investigator, International Ultraviolet Explorer satellite.

temperature-gravity variations of the line profile across the

gravity-darkened stellar disk should be included but also tin

principle) the effects of stellar winds, nonradial pulsations, and

(possibly) differential rotation. All of the latter may affect the

line profiles, but cannot easily be taken into account at this

time.

As Huxhings (1976b) and Sonneborn and Collins 11977)

have pointed out, if rotational broadening is indeed a measurabte function of wavelength, it becomes possible in principle to

gain knowledge of global properties of rotating stars and, specifically, to determine v and i separately. Indeed. such determinations have been attempted in recent years by Hutchings and

Stoeckley (1977), Hutchings. Nemec, and Cassidy 11979), and

Ruusalepp (1982). Are these results meaningful, in view of the

uncertain assumptions and difficulties that have been pointed

out by :Marlborough (1982)?

In the hope of clarifying the situation somewhat, we have

calculated new theoretical line profiles based on line-blanketed

rotating model atmospheres, which include shape distortion

and gravity darkening, plus the latest atomic data, and compared our results with ultraviolet line profiles in a number of

later B type and A type stars observed with the I UE satellite.

11. THE THEORETICAL ULTRAVIOLET LINE PROFILES

Atomic absorption lines chosen for rotational velocity

analysis should be (1) as free as possible from broadening or

profile-distorting mechanisms other than rotation and (2) sufficiently strong and isolated from other lines to allow accurate

profiles to be determined. These conditions present difficulties

in the ultraviolet, where the lines tend to be crowded together

and the strongest lines are inevitably resonance lines. Resonance line profiles can indeed often be measured with high

precision but are also likely to be distorted by stellar winds, in

normal B type as well as in Be stars, even in the later B types

(cf. Marlborough 1982; Snow 1982; Slettebak and Carpenter

1983).

We have chosen two moderately strong nonresonance lines.

which we assume to be photospheric and broadened essentially

by rotation in our sample of later B type and A type stars. For

the spectral range B5-B8, we use the feature near 1299 .A which

is produced by the Si III lines at 1298.960 and 1298.891

A. These lines have their lower states at about 6.5 eV and their

741

DIP a
Q)

742 CARPENTER, SLETTEBAK, AND SONNEBORN Vol. 286

higher energy levels near 16.1 eV. The neighboring Si in lines at

1297 and 1301 A were included in our calculations but were

judged to be too weak in the observed spectra to be useful in

the analysis. Finding suitable lines for the A type stars is more

difficult, since they have little energy shortward of 1500 A and

the spectra are crowded with lines. The best compromise we

could make was the 2755.733 A line of multiples UV62 in the

spectrum of Fe it. This line is fairly strong and relatively isolated, and is produced from a lower energy level at 0.98 eV to a

higher level at 5.46 eV.

While not resonance lines, both the Si in X1299 and Fe it

X2756 lines have lower levels that do not connect strongly to

the ground state. In this sense, these are metastable levels, and

absorption lines arising from them will be enhanced under

low-density conditions, which will overpopulate these leve-Is.

Our lines, therefore, should not be used for determining rotational v ::_ -ties of shell stars, since they would yield estimates

too low relative to the values of u sin i measured from lines that

originate solely in the photosphere.

Synthetic spectra of regions 6 A wide, approximately centered on each of the two lines, were computed by using a

modified version of the ATLAS6 (Kurucz 1979) model atmosphere code, the global intensity integration routines of Collins

and Harrington (1966), and previously computed model atmospheres described in Slettebak, Kuzma, and Collins (1980). The

physical parameters for these line-blanketed model atmospheres, which account for the shape distortion and gravity

darkening caused by rapid rotation (resulting in the variation

of the photospheric continuum brightness with latitude), are

given in the latter reference. Data for the atomic lines included

in each of these synthetic spectra were obtained by screening

the million-line list of Kurucz (1979) to obtain data for the

strongest lines in each wavelength region under the physical

conditions expected in each model. Since data for the Si in

(UV4) lines are not in the Kurucz (1979) compilation, their

wavelengths, gf values, and lower energy levels were taken

from Wiese, Smith, and Miles (1969). The Kurucz data were

screened by cemputing approximate line-center opacities corresponding to the temperatures, pressures, and densities near

Rosseland optical depth r - J in the midlatitude atmospheres

( --45°) of each stellar model. The 100-200 strongest lines were

segregated by retaining for each model only those lines with

central opacities greater than 9.0 cm' g - 1 . The Fe It line and all

the background lines were computed using the wavelengths,

energy levels, and gf-values from the Kurucz tape, and the

broadening approximations described by Kurucz (1979) and

Kurucz and Furenlid (1979). The approximate damping constants are defined as follows:

r(stark) : i x elf

r(radiative) = 2

2. 223 x 1015 (^ in A)

A

f (van der Waals) = 2.8393081

10 -

"(NH , + 0.42NH , I)T a IR2

where

_ = 2 4 -2 R — .Sn,rr Z.rr

and

2

/'

13.5952.,r

norr 

= (,E — E,., Inel + 0.1 )

with NHI and NHe1 the number densities of H 1 and He I.

respectively, n,rr the effective quantum number, and Z,,, the

ionization stage. The ratio A of the damping to the Doppler

width is defined by

A= 

r(radiative) + r(van der Waals) + ['(Stark)

4rrAVD

Collisional broadening of the Si in (UV4) lines was approximated using the semiclassical results of Sahal-Brechot and

Segre (1971), for which

r, = Q I N, TQ:

The values of Q, = 3.5 x 10' and Q2 = -0.3 used in this

equation were taken from Kamp (1976). The radiation

damping constants, ('(radiative), for Si to (UV4) were approximated by the upper-state Einstein A-values. A microturbulent

velocity of 2 km s - ` was assumed for all lines in all of the

models. These spectra were calculated with a wavelength

spacing A. - 0.05 A, with the central wavelength of each line

rounded to the nearest grid point to ensure that each line is

represented in the spectrum.

Since rotational broadening dominates the profiles of the

lines being considered for the v sin i values of interest, our

results should not be strongly dependent on 'he above assumptions.

Table 1 lists the rotational parameters of the models for

which line profiles were computed, including values of the fractional angular velocity w = co/co, (where w, is the critical

angular velocity for which the centrifugal force at the star's

equator balances the gravitational force), the equatorial rotational velocity v, and the inclination angle i between the rotation axis and the line of sight. Si in line profiles were calculated

for the B5-AO models, and Fe it line profiles for the AO-A7

models. Figures 1-3 show selected examples of these theoretical line profiles. For each profile, the full width of the line at

half-maximum depth IFWHM) was measured, with the continuum taken to be at a residual intensity of 1.00, and assumed to

be an indicator of the rotational velocity of that model. The

signal-to-noise ratio in the observed /UE line profiles does not

justify the use of additional line-width parameters in the

analysis, in our opinion. Calibration curves with the measured

FWHMs as a function of v sin i were then plotted for each line

and model. Figure 4 shows such calibration curves, as an

example, for the Fe it . 2756 line in A5 models.

TABLE I

ROTATIONAL PARAMETERS OF MODELS FOR WHICH LINE PROFILES WERE

COMPUTED

u sill i (knl S - ') AT EACH SPECTRAL TYPE

W i B5 B7 AO A2 A5 A7

0 ...... 0 0 0 0 0 0

0.5 .. ..... W 0 0 0 0 0

30 70 64 66 65 65

45 98 91 93 92 92

60 120 Ill 114 112 113

90 139 128 132 129 130

0.9... ..... 0 0 0 0 0 0 0

30 145 136 134 138 135 135

45 205 193 189 195 191 191

60 251 X36 232 239 234 234

90 290 273 268 276 210 270
1.00

85 MODEL (w n 0.9)

SIX (UV4)

1296.726 1296'910 1301.146 1 91.691 1

1.00

AO MODEL (w • 0.5)

F* It (62) 2755.733 A

1

0.50

Y

F

fA

2

W

H

2

J

a

O

HW

M

Y

H

N

2

W

H

2

J

a

0

to

W

cc

0

1.00

0.50

0•

v sin i - 0 km/s

i • 30•

v sin i - 145 km /s

- 0•

v sin i 0 km /s

30 .

V sin - 64 km /s

0.50

0

1.00

0.50 r n

.,

No. 2, 1984 ROTATIONAL VELOCITIES OF LATER B AND A STARS 7,11

1.00

. SW

V sin i - 290 km /s

0.50

1296 1296 100 1302

WAVELENGTH (A)

FIG. I.—Theoretical line profiles of Si to multiplet UV4 lines in the wavelength range 1296-1302 A, computed for 85 models with w s 0.9 and three

values of the inclination i.

Ill. THE PROGRAM STARS

a) Spectral Types and Rotational Velociticsfrom Visual

Spectra

Since we already had I UE high-resolution spectra of a

number of rapidly rotating standard stars of late B type and A

type, we decided to concentrate on this spectral-type range.

Our requirements were that the program stars have welldetermined spectral types and rotational velocities (from visual

spectra) and that high-resolution 1 UE spectra exist. We also

tried to include stars with a range of rotational velocities in our

sample, although highest priority was given to rapid rotators,

where gravity-darkening effects are predicted to be the largest.

Our 17 program stars are listed in Table 2. The spectral types

and rotational velocities are from Slettebak et al. (1975), Slettebak (1982), and Slettebak and Carpenter (1983), or references

therein. Details regarding how the visual v sin i values were

determined, as well as error discussions, may be found in the

aforementioned references. We estimate uncertainties of

10-15% of the values listed in Table 2.

b) The Observed Ultraviolet Line Profiles

and Rotational Velocities

Table 3 lists the ultraviolet spectroscopic information for

our program stars. All were observed with the lUE satellite in

the high-dispersion mode, through the large aperture. We

obtained eight of the spectra directly during the fifth and sixth

observing episodes of the I UE (1982-1983) and the remaining

nine from the /UE archives of the National Space Science Data

Center (NSSDC). The star name and either the shortwavelength prime (SWP) or the long-wavelength redundant

100

• 90'

V sin i • 126 km /s

0.50

2753 2755 2757 2759

WAVELENGTH (A)

FIG. 2.—Theoretical line profiles of the Fe n multiplet UV62 line at

2755.733 A. computed for AO models with w •- 0.5 and three values of the

inclination i.

(LWR) camera image number are given in the first three

columns of the table. Line widths (FWHM) were measured on

the 1 UE spectra, corrected for the finite resolution of the I UE

high-resolution camera, and compared with the line widths of

the appropriate theoretical line profiles to obtain rotational

velocities. Here an ambiguity enters, which is shown in Figure

4 and was discussed earlier by Collins (1974) and by Slettebak

et al. (1975). These, as well as our present computations, show

TABLE 2

THE PROGRAM STARS

Star HD Spectral Type

Visual

v sin i

Ikm s')

p Aur ...... .. 34759 B5 V 90

y ' Aqr _ .. 219688 B5 V 280

U Psc ........... 217891 B5 Ve 100

B CrB ..... - ... 138749 B6 III(e) 320

m And .......... 6811 86.5 IIle 80

z Leo .......... 87901 B7 V 280

?f Tau .......... 23630 B7 Ille 140

9 Cyg 8 ........ 183914 B8 Ve 250

And ........... 222173 B8 V 40

I Peg . ......... 218045 B9.5 III 120

o Cyg .......... 186882 B9.5 111 140

B Car ........ .. 80007 A I IV 125

: PsA .. ....... 216956 A3 V 85

80 UMa ........ 116842 A5 V 210

6 Cas ........... 8358 A5 III-IV 100

2 Oph .......... 159561 A5 111 215

s Aql ........... 187642 A7 V 225
, = 30•

v tint • 135 km/s

744 CARPENTER, SLETTEBAK, AND SONNEBORN Vol. 286

A5 MODEL (w n 0.9)

100 Fog

(62)

2755.733 E

•

0.50 A

v un % - 0km/s

H

2

WtZ 0

J 1.00

Q

O

N

W

tL

0.50

1.00

i v 90•

v fin, • 270 km/s

0.50

2753 2755 2757 2759

•

WAVELENGTH (A)

FIG. 3.-Theoretical line profiles of the Fe a multiplet UV62 line at

2755.733 .4, computed for A5 models with w - 0.9 and three values of the

inclination ,.

that the half-intensity width does not uniquely specify a rotational velocity for large values of w. We therefore took an

average v sin i where both w = 0.5 and w = 0.9 models were

applicable. This ambiguity introduces an uncertainty of

15-200 in our rotational velocities, listed in Table 3.

FWHM(A) F*II 2756

3.0 i

a^

i

i s^

2.5

2.0

1.5

1.0

A5 MOOELS

o-o W 0.5

0.5

"-• W = 0.9

0 -

0 50 100 150 200 250 300

v sin i (km/sec)

FIG. 4.-Line widths (FWHM) of theoretical line protnes of the Fe n /.j756

line in A5 model rotating stars. plotted against the'r rotational velocities, r sin

i. The circles and crosses represent computed values for i = 0 . 30`. 45 60 .

and 90°; the dashed curves approximate the mean relations. Error bars of

= 10%. which should represent maximum errors in view of uncertain continuum placements, are shown on the FWHM.

IV. RESULTS AND DISCUSSION

Figures 5 and 6 show the ultraviolet rotational velocities

derived in this study from the Si tit ;, 1299 and Fe a 1".2756 lines,

respectively, plotted against the visu-ily determined r sin i

values for the 135-88 and 88-A7 program stars. The agreement

between the values of v sin i is quite good in both cases, and we

conclude that our models reasonably represent the observed

line widths.

Sonneborn and Collins (1977) introduced the dimensionless

parameter

Q = IFWHMii.) I ,(FWHM /'.),

TABLE 3

ULTRAVIOLET DATA FOR THE PROGRAM STARS

ST I.R

CAMERA IMAGE No.

SWP LWR

Si III .11299

FWHM 1"A)

Fe a A2756

FWHM (A)

UV

v sin i

(km s')

p Aur ..... _ . _ .. 10389 .. 0.80 .. 70

w 1 Aqr ......... 10385 1.52 ... 255

0 Psc .......... 10386 0.86 I 85

15512 _. 0.84

0 CrB .......... 14431 163 335

16288 1.69 _.

o And ........ 10387 0.77 65

: Leo........... 10379 1.55 ... 280

q Tau ......... 10378 1.19 185

!l C,vg B ....... 5369 1.41 235

And .... .. 10376 0.82 85

9057 1.31 8o

X Peg ..... .... 2386 _ . 2.23 145

b Cyg _..... 3041 2.11 140

0 Car 2970 1.69 105

z PsA . ...... 11022 1.81 110

80 UMa .. ..... 11712 _ 2.57 190

o Cas .......... 11214 ... 1.92 75

7 Oph ......... 5925 2.72 220

s Aql ........... 2449 3.06 ) 190 3012 2.79 1

f
No. 2, 1984 ROTATIONAL VELOCITIES OF LATER B AND A STARS 7.4;

UV — v sin i (km/sec)

/

300

250

200

150

100

Sim 1299

50

50 100 150 200 250 300

VISUAL - v sin i (km/sec)

FIG. 5.—Visually determined rotational velocities ( u sin i) plotted against

ultraviolet u sin i values ( from Si to .11299 lire profiles) for the B5-88 program

stars. Error bars of ± 15% have been placed on each point. he dashed line is

a 45` line representing perfect agreement between the visual and ultraviolet u

sin i values.

as a measure of the difference in rotational line broadening as a

function of wavelength. We can compute Q from our ultraviolet computations in this paper and the visual line profiles calculated by Slettebak, Kuzma, and Collins (1980). This gives

values of Q between 1.2 and 1 . 3 for B5 models with w = 0.9 and

i = 90°, and agrees quite well with observed Q-values for the

rapidly rotating B5-117 stars in this study obtained from IUE

spectra compared with He t A4471 lLie widths for the same

stars in Slettebak et al. (1975). This confirms that there is an

UV- v sin i (km/sec)

250

200

150

loo

-^- 88-A7

50 / Fe 11 2756

50 100 150 200 250

VISUAL - v sin i (km/sec)

FIG. 6.—Visually determined rotational velocities to sip 4 plotted against

ultraviolet v stn i values ( from Fe it ~2756 line profiles) for th- 88-A7 program

stars. Error bars of - 15^s, have been placed on each point. The dashed line is

a 45 line representing perfect agreement between the visual and ultraviolet v

sin i vab­ss.

effect of wavelength dependence on rotational line broadening

in the B type stars, that it can be understood in terms of

gravity-darkened models, but that it is probably not as lar ge as

previously thought.

Our Q-values for the A type stars in this study, on the other

hand, are near unity, even for the most rapidly rotating objects.

There are probably several reasons for this. (1) Even the most

rapidly rotating A type stars appear to have fractional angular

velocities that are not sufficiently close to unity to produce

striking gravity-darkening effects. (2) Since the temperature difference between pole and equator in a rapidly rotating A type

star is considerably less than in a rapidly rotating B t y pe star.

the gravity-darkening effects would be expected to be much

smaller in any case. (3) The Fe a A2756 line, which we chose for

analysis in the A type spectra, is not so widely separated in

wavelength from the Mg it ;.4481 line usually chosen for visual

u sin i determinations as to show a large wavelengthdependence effect.

Before our conclusions are taken too seriously by the reader,

however, we feel it is necessary to comment on the accuracy of

our observed and theoretical line profiles, as well as on some

ambiguities in our analysis:

1. Placing the continuum for the observed line profiles is

extremely uncertain. The high-resolution IUE spectra are

obtained with an echelle spectrograph, and establishing continuity between echeile orders is difficult and not always successful. The problem of locating the continuum becomes worse

when the line profiles to be measured are broad and other

strong lines are in the vicinity. Also, for broad lines, a small

error in continuum placement can result in a relatively large

error in measured line width and, therefore, rotational velocity.

2. There is also some uncertainty in the location of the continuum for the theoretical line profiles, because of the restricted

(by limitations on available computing time) wavelength extent

of the spectra computations.

3. Figure 4 shows two problems introduced in using the

FWHM versus u sin i calibration curves. The first, which we

have already discussed, is that there is an essential ambi guity in

using line widths to estimate rotational velocities for rapidly

rotating stars, as is shown by the difference between the

w = 0.5 and w = 0.9 curves. A second source of uncertainty

arises from the flattening out of the FWHM versus c sin i

curves for large values of v sin i. A small difference in the

measured line width thus leads to a large change in u sin i.

The above discussion suggests that Figures 5 and 6 be

regarded with some caution. The rather good correlation

i4tween the visually determined and ultraviolet c sin i values

shown there is probably somewhat fortuitous. On the other

hand, we feel that our anal ysis is sufficientl y accurate to reveal

striking differences, and these do not seem to be present.

Improvements in both theory and observation could be

made in future work. As we have already stated, we feel that

our theoretical treatment is sufficiently accurate to deal with

this problem in view of the fact that rotational broadening

,.,iminates the line profiles. It would be advantageous.

however, to compute synthe!ic spectra covering a larger wave-

'ength range than was possible in this paper, to define the

theoretical continuum more precisely. On the observational

side, we anticipate that the quality of the line profiles obtained

with the high-resolution spectrograph on the Space Telescope

would permit a comparison of the entire line profile, rather

than only the F`N HMs used in this investigation, and therefore

would lead to more accurate results. The choice of ;uitable

s

A
U

746 CARPENTER, SLETTEBAK, AND SONNEBORN

lines for analysis presents real problems, since all strong,

unblended ultraviolet lines seem to arise either from ground

states or from metastable levels. We have chosen what we

consider to be the best lines for our purpose, but higher

resolution, less noisy data may permit the use of more suitable

lines.

The many uncertainties inherent in analyses involving large

rotational line broadening also suggest to us that attempts to

separate u and i from line profiles must be carried out with

great caution, and, indeed, it is not clear to us that iuch a

separation is even feasible.

We can summarize our results as follows. Real differences in

widths for ultraviolet versus visual line profiles in rapidly rotating B type stars exist, but these are not as large as previously

reported in the literature. This discrepancy may be due in part

to the earlier use of resonance and metastable lines that have

substantial nonphotospheric contributions from winds and/or

circumstellar envelopes. The good agreement we obtain

between ultraviolet and visually determined rotational velocities suggests that our shape-distorted, gravity-darkened

models are reasonable and may bear some resemblance to the

real world.

We acknowledge very helpful conversations with George

Collins. The 1UE team at Goddard Space Flight Center was

most cooperative, as was Wayne Warren of the National Space

Science Data Center at Goddard. We thank also the Ohio

State University Instruction and Research Computer Center

for providing computer time and facilities, and NASA for

support via grant NAG 5-52.

REFERENCES

Collins, G. W., 11. 1974, Ap. J., 191, 157.

Collins, G. W. 11, and Hamngton, J. P. 1966, Ap. J., 146,151

Heap, S. R. 1975, Phil. Trans. Roy. Soc. London, A, 279, 371.

1976, in 1AU Symposium 70, Be and Shell Stars, ed. A. Slettebak

(Dordrecht. Reidell, p. 165.

. 1977, Ap. J. (Letters), 218, L 17.

Hutchings, J. B. 1976x, in lAU Symposium 70, Be and Shell Stars, ed. A.

Slettebak (Dordrecht: ReidelL P. 450.

. 1976b, Pub. A.S.P., 88, 5.

Hutchings, J. B.. Nemec, J. M., and Cassidy, J. 1979, Pub. A.S.P., 91.313.

Hutchings, J. B., and Stoeckley, T. R. 1977, Pub. A.S.P., 89, l9.

Kamp. L. W. 1976. Statistical Equilibrium Calculations for Silicon in EarlyType Model Stellar Atmospheres(NASA TR R-455).

Kurucz, R. L. 1979, Ap. J. Suppl., 40, 1.

Kurucz, R. L., and Furenlid.1. 1979, Smithsonian Ap. Obs. Spec. Rept., No. 387.

Marlborough, J. M. 1982, in lAU Symposium 98, Be Stars, ed. M. Jaschek and

H.-G. Groth IDordrecht: Reidel), p. 361.

Morton, D. C., Jenkins, E. B_ Matilsky, T. A., and York, D. G. 1972 Ap. J.,

177, 219.

Ruusalepp, M. 1982, in 1AU Symposium 98, Be Stars, ed. M. Jasehek and H.-G.

Groth (Dordrecht: Reidel) p. 303.

Sahal-Brechot, S., and Segre, E. R. A. 1971, Asrr. Ap., 13.161.

Slettebak, A. 1982, Ap. J. Suppl., 50, 55.

Slettebak, A., and Carpenter, K. G. 1983, Ap. J. Suppl., 53,869.

Slettebak. A.. Collins, G. W., 11. Boyce, P. B.. White. N. M., and Parkinson. T.

D. 1975, .9p. J. Suppl., 29,137.

Slettebak, A., Kuzma, T. J., and Collins, G. W.. 11.1980. Ap. J., 242, 171.

Snow, T. P., Jr. 1982, in lAU Symposium 98, Be Stars, ed. M. Jaschek and

H.-G. Groth (Dordrecht: Reidel), pp. 377.

Sonneborn, G. H., and Collins, G. W., II. 1977, Ap. J., 213.787.

Wiese, W. L., Smith, M. W., and Miles, B. M. 1969, Atomic Transition Probabilities, Vol. 2 (NSRDS-NBS 22).

KENNETH G. CARPENTER: JILA, University of Colorado, Boulder, CO 80309

ARNE SLETTEBAK ' Perkins Observatory, P.O. Box 449, Delaware, OH 43015

GEORGE SONNEBORN: 1UE Observatory, Code 685.9/CSC, NASA/Goddard Space Flight Center, Greenbelt, MD 20771

1
a) Ultraviolet spectra of some brighL s.' Lambda Bootis stars

b) A. Slettebak1

c) Perkins Observatory, Ohio State and Ohio Wesleyan Universities

Delaware, Ohio 43015, U.S.A.

d) ULTRAVIOLET SPECTRA OF LAMDA BOOTIS STARS

e) Perkins Observatory

P. 0. Box 449

Delaware, Ohio 43015, J.S.A.

f) Stellar Atmospheres

g) Supplement Series

1 Guest Investigator: International Ultraviolet Explorer Satellite
0

-2-

Summary. - Strengths of ultr • . .Lolet lines from IUE spectra of six bright

X Bootis or suspected a Bootis stars are compared with line strengths

in standard "normal" stars of similar spectral types. In five of the six

stars (the sixth, E Sgr, aplears to be normal in the ultraviolet) the lines of

carbon and nitrogen are stronger than in the standard stars, while the lines

of magnesium, aluminum, and (to a lesser extent) iron are weaker. These

abnormal line strengths, which seem not to be due to luminosity effects, might

be interpreted either in terms of diffusion p rocesses or by nucleosynthesis.

The line profiles of ultraviolet resonance lines do not show the asymmetries

associated with mass loss in the hotter stars, but the low mass-loss rates

suggested in the diffusion theory are probably not excluded by the observations. On the other hand, the observed strength of ultraviolet carbon and

nitrogen lines is suggestive of nucleosynthetic processes, and the possibility exists that the X Boo stars are evolved objects in which nuclearlyprocessed CNO elements have somehow found their way to the surfaces.

Key Words: UV spectra -- Lambda Bootis stars -- peculiar A stars --

stellar atmospheres
rV

-3-

I. Introduction

Morgan at al. (1943) described the spectrum of X Bootis, the prototype of this

class of stars, as follows: "The spectral type of 'A Boo is near A0, as far as

can be determined. The spectral lines, while not unusually broad, are veil

weak, so that the only features easily visible are a weak K-line and the

Balmer series of hydrogen". Additional members were proposed by Andersen and

Nordstrom (1977), Graham and Slettebak (1973), Parenago (1958), Slettebak

(1952, 1954, 1975), and Slettebak et al. (1968), but it was obvious that the

class was rather ill defined. Hauck and Slettebak (1983) therefore proposed

the following spectroscopic definition: % Boo stars are A-F type stars with

metallic lines which are too weak for their spectral types, when the latter

are determined from the ratio of their K-line to Balmer-line strengths. They

are distinguished from weak-lined Population II (horizontal branch) stars by

the fact that they have normal space velocities and moderately large rotational velocities." Based on this definition, they listed nine stars which

appear to be a Boo stars. Meanwhile, Abt (1984), suggesting that "... an easy

way to discover a Bootis stars is to search for stars near AO that have weak

(Mg II) 4481", presented a list of 26 stars near AO selected from the Bright

Star Catalogue which have weak Mg II 4481, and which he provisionally refers

to as a Bootis stars.

Abundance analyses of some :)f. ,:;:,e brighter a Boo stars have been carried

out by a number of investigators. Burbidge and Eurbidge (1956), using curveof-growth methods, found X Boo to have only about 10 percent of the "normal"

abundance of Fe and only about 3 percent of the to abundance of Ca, with

9-alle y underabundance ratios for 29 Cyg. Kodaira (1967) found % Boo and
-4-

rr l Ori to have normal oxygen abundance, in contrast to the underabundance of

magnesium by a factor of more than 5." For 29 Cyg, on the other hand, he

found the magnesium abundance to be "rather normal". Baschek and Searle

(1969) carried out a curve-of-growth abundance analysis of a number of objects

which had been classified as a Boo stars, and found that "only a Boo, 29 Cyg,

and Tr l Ori form a distinct group from the composition point of view. The composition characteristics of the A Boo group are (1) a deficiency of about a

factor 3 in the iron-group elements; (2) nearly normal oxygen abundance; 3)

deficiencies by Mg and Ca that are at least as large as those of the irongroup elements and sometimes Pre larger."

Baschek and Searle (1969) agreed with Sargent's (1965) suggestion that

the a Boo stars are related to the peculiar A stars, and stated further that

"... they form a separate type of Ap star, and their surface composition does

not reflect that of the material form which they formed." More recently,

Michard et al. (1983) have suggested that the a Boo stars may be related to

the A-type metallic-line (Am) stars. They propose that the diffusion process

(Michaud 1970) plus a mass-loss rate of the order of 10 -13 Me /yr could lead to

the % Boo phenomenon.

An ultraviolet study of nine a Boo or suspected a Boo-type stars was

carried out by Baschek et al. (1984) using IUE spectra. Of five stars listed

by Hauck and Slettebak (1983) as Boo stars based on optical spectra, they

found that four (a Boo, rr l Ori, HR 4831, anI HD 105058) also show ', Boo

characteristics in the ultraviolet, while the fifth (29 Cyg) shows only

marginal agreement. A study of high-resolution IUE spectra of 'k Boo itself
-5-

VW J

led them to conclude that "the line strengths indicate that carbon (and

perhaps also nitrogen) may be normal or even enhanced in a Boo in contrast to

the underabundance of the heavier elements."

The work of Baschek at al. (1984) was not known to me when I also obtained IUE spectra of some of the brighter a Boo stars in October 1983 and January

1984. My goal was to compare high-resolution IUE spectra of a Boo stars with

spectra of standard "normal" stars of similar spectral type, to determine

which chemical elements have stronger, weaker, or normal lines in the a Boo

stars. This paper presents these results, based on independent IUE observations. No attempt is made to analyze the results or derive abundances,

since the data do not warrant a refired analysis, in my opinion. The lines of

certain elements do appear to be strong and others weak, relative to their

appearance in st;-%dard stars, however, and these will be pointed out.

II. Observations

Table la lists the six a Boo or suspected a oo-type stars for which highresolution IUE spectra were obtained, including the image numbers. Table lb

gives the same information for 14 standard stars with normal optical spectra,

of spectral types B8 to A5. Spectral types and rotational velocities for the

peculiar stars are referenced in the paper by Hauck and Slettebak (1983),

except for 15 And (see the ;Vote to Table la), while those for the standard

stars are from Slettebak et al. (1975) or references therein.

Identifying representative spectrum lines of various chemical elements

which are unblended and at the same time strong enough to perm't accurate

r
-6-

equivalent width measurements is a near impossibility for ultraviolet spectra

of hot stars. Furthermore, many or most such candidate lines are resonance

lines, which introduces complications due to interstellar and/or chromospheric

contributions. Recognizing these difficulties, Table 2 lists the lines (many

of them blended) which seemed least objectionable for the comparison of line

strengths, which is the purpose of this invests;stion.

ME ;alcomp tracings of the spectra of al. LhP ,:cogram stars and all the

standard stars but two (18 Tau and a UMd) were available with a paper scale of

2 R/inch. Equivalent widths were measured on these tracings, using a planimeter. In addition, equivalent widths of certain lines in the spectra of

18 Tau and a UMa, as well as in some of the program stars, were measured using

a reduction program at the WE Regirnal Data Analysis Facility in Boulder,

Colorado. Because a considerable range of continuum is necessary on either

side of each spectrum line to be measured in order to obtain meaningful

equivalent widths, the Boulder reduction program was found, unfortunately, to

be useful only for certain lines.

III. Results

Baschek et al. (1984) have pointed out the difficulties in reaching

quantitative conclusions regarding the abundance of carbon in the atmosphere

of N Boo in view of saturation effects in the lines (making them sensitive to

microturbulent velocities and damping constants) and the possibility of

differential non-LTE effects on the line strengths. There are also a number

of difficulties on the observational side: (1) As has already been mentioned, ultraviolet spectra of the hotter stars are crowded with lines and it is

t
-7-

extremely difficult to find lines which are unblended and at the same time

strong enough to permit accurate equivalent-width measurements; (2) the

tracings utilized in this study often had poor flux corrections at the ends of

the echelle orders ("ripple" corrections), which made continuum placement for

lines at these wavelengths very difficult; (3) unfortunately, a number of my

tracings were very noisy, which affected both the line profiles and continua.

For all of the above reasons, only line strengths of the program stars

versus standard stars are presented in this paper, witA no attempt at

analysis. My hope is that, by measuring the lines . in the a Boo stars and in

the standard stars of similar spectral types in the identical manner, meaningful comparisons can be made.

The results are shown in Figures 1-24, one for each of the spectrum lines

listed in Tables 2a acid 2b. In each case, equivalent width W(R) it plotted

against spectral type for both the program and the standard stars.

An inspection of Figures 1-24 s!iows, first, a considerable scatter in the

points, making it difficult to decide whether a given line in the program stars

is weaker, stronger, or comparable to the line in the standard stars. In a

number of cases, the comparison shows a definite difference, however:

1. Figures 3 and, especially 7, suggest that the lines of carbon are

decidedly stronger in the A Boo stars (with the exception of E Sgr) than in

the standard stars of comparable spectral type. This supports the conclusion

of Baschek et al. (1984) that carbon may be normal or even enhanced in ^ Boo.

2. Figures 4, 5, 10, and 11 suggest that the lines of nitrogen are also decidedly stronger in the ^ Boo stars (again, with the exception of E Sgr) than

k^

C-
-8-

in the standard stars of comparable spectral type. Baschek et al. (1984) had

also suggested that the lines of nitrogen in .^ Boo "perhaps" may be normal or

even enhanced. The N I (multiplet W 4) lines at 1492.63, 1492.82, and

1494.68 R are, unfortunately, blended with the C I (multiplet 64.06) 1494.74

and C I (multiplet 64.08) 1492.74 R lines in an unknown way, but Figures 10

and 11 show that the N I (multiplet UV 9) lines at 1743 and 1743 R are also

strong relative to standard stars.

3. Figures 21-24 show that the lines of magnesium are decidedly weaker in the

X Boo stars (again, with the exception of a Sgr) than in the standard stars of

comparable spectral type. This might be expected since the weakness of Mg II

4481 is one of the criteria for identifying a Boo stars. Burbidge and

Burbidge (1956), Kodaira (1967), and Baschek and Searle (1969) all found

magnesium to be underabundant in the ^ Boo stars.

4. Figure 8 shows that the A2 II 1671 resonance line is considerably weaker

in the ^ Boo stars (again, with the exception of E Sgr) than in the standard

stars of comparable spectral type. It is interesting to note that Sadakane et

al. (1983) found the resonance lines of both AZ II and AZ III to be abnormally

weak in many B-type chemically peculiar stars.

5. Figures 14, 16, 17, 19, and 20 show the lines of Fe II to be definitely

weaker or slightly weaker in the X Boo stars (again, with the exception of

E Sgr) than in the standard stars of comparable spectral type. The effect is

not as pronounced as for the Mg II and AZ II lines, but seems to be real.

6. Figure 9 shows the Ni II 1742 line to appear slightly weaker, on the

average, in the 1 Boo stars (again, with the exception of E Sgr) than in the
C9 71

-9-

standard stars of comparable spectral type. This effect may not be real, in

view of the uncertainty of the equivalent width measurements.

7. No definite conclusions can be drawn about the strengths of the lines of

silicon (Figure 6), sulfur (Figure 12), or manganese (Figures 15 and 18), in

view of the uncertain measurements, except to suggest that the lines seem

neither to be conspicuously strong nor conspicuously weak relative to the same

lines in the standard stars.

A number of other conclusions may be drawn from the data:

The bright star eSgr seems not to be a a Boo star on the basis of the

strengths of its ultraviolet lines. In addition to the results already

mentioned, Figures 1, 2, and 13 show that lines of Si III, 0 I, and AZ III,

respectively, in ESgr appear to have normal strength: relative to standard

stars of similar spectral type. On the other hand, the optical spectrum of

E Sgr is peculiar in that the Mg II 4481 line is broad and rather weak while

the Balmer absorption cores, Fe II 4233, and the Ca II K-line are considerably narrower, as was pointed out earlier (Slettebak 1975). The spectrum

could be composite, but again there is no evidence for abnormal line intensities in the ultraviolet. Hauck and Slettebak (1983) have pointed out

that the colors of E Sgr are unusual and that its position in both a d vs.

(B2-Vl) diagram and a c vs. (b-y) diagram place it in the area of the

horizontal-branch A-type stars. This object does indeed deserve further

study.

Abt's (1984) suggestion that 15 And is a 1 Boo star on the basis of its

weak Mg II 4481 line seems to be confirmed by the IUE data. The star shares

the strong C I and N I lines, the weak Mg II and AZ II lines, and the slightly
-10-

weak Fe II lines shown by the ether A Boo stars studied in this paper. On the

other hand, Hauck (1985) decided against membership of 15 And in the a Boo

group on the basis of its colors'in the Geneva photometric system. Hauck's

criteria for membership are severe, however, and he states that "some stars

(of the objects investigated by him) could marginally be members of the group

.... with only one clear exclusion (an obvious supergiant)".

In view of Michaud et al.'s (1983) suggestion that the a Boo stars may be

losing mass, the ultraviolet resonance line profiles were examined for

asymmetries in all of the program stars. No asymmetries were found. On the

other hand, Michaud et al.'s suggested mass-loss rate of 10 -13 Mo /yr is not

much larger than the mass loss due to the solar wind -- under these circumstances there would probably be no measurable effect on the line profiles.

IV. Conclusions

1. Five of the six A Boo stars investigated in this study have similar ultraviolet spectral characteristics; the lines of carbon and nitrogen are

stronger than in standard stars of similar spectral type, while those of

magnesium, aluminum, and (to a lesser extent), iron are weaker. Nickel may

also be weakened.

2. The sixth star investigated. a Sgr, while showing peculiarities in its

optical spectrum and colors, looks normal for its spectral type in the ultraviolet.

3. The data presented in Figures 1-24 show no striking differences in line

strengths between main-sequence and giant stars, suggesting that the behavior
r-a11

r

i

K ' ^

V: -11-

of the ultraviolet lines in the .^ Boo stars is not due to luminosity effects.

The abnormal line strengths in the ultraviolet spectra of the a Boo stars

might then be interpreted either'in terms of diffusion processes, as suggested

by Michaud et al. (1983), or by nucleosynthesis. The line profiles of ultraviolet resonance lines do not show the asymmetries associated with mass loss

in the hotter stars, but the low mass-loss rates suggested by Michaud et al.

are probably not excluded by the observations. On the other hand, the

observed strength of ultraviolet carbon and nitrogen lines is suggestive of

nucleosynthetic processes (Kodaira 1967 and B;schek and Searle 1969 had

earlier found that the oxygen abundance is normal in a Boo stars). The

possibility thus exists that the a Boo stars are evolved objects in which

nuclearly-processed CNO elements have somehow found their way to the surfaces.

Acknowledgments. I thank Charles Cowley for restimulating my interest in the

and &vie LucK

r Boo stars, and am grateful also to him and to George Collins Afor valuable

discussions. The IUE team at the Goddard Space Flight Center was extremely

helpful, as was Wayne Warren at the National Space Science Data center in

sending archival material to me. My thanks go also to Edward Brugel and Terry

Armitage of the IUE Regional Data Analysis Facility in Boulder, Colorado for

their help and cooperation during my visit there, and finally, to NASA for

support via grant NAG 5-52.

i

^ :. I - _. J^
10'

-12-

References

Abt, H. A.: 1984, in The MK Process and Stellar Classification, ed. R. F.

Garrison, David Dunlap Obs.•Publ., p. 340.

Andersen, J., Nordstr8m, R : 1977, Astron. Astrophys. Suppl..29, 309.

Bashek, B., Heck, A., Jaschek, C., Jaschek, M., Oppen, J., Scholz, M.,

Wehrse, R.: 1984, Astron. Astrophys. 131, 378.

Baschek, B., Searle, L.: 1969, Astrophys. J. 155, 537.

Burbidge, E. M., Burbidge, G. R., 1956, Astrophys. J. 124, 116.

Graham, J. A., Slettebak, A.: 1973, Astron. J. 78, 295.

Hauck, B.: 1985, Astron. Astrophys. (in press).

Hauck, B., Mermilliod, M.: 1980, Astron. Astrophys. Suppl. 40, 1.

Hauck, B., Slettebak, A.: 1983, Astron. Astrophys. 127, 231.

Kodaira, K.: 1967, Publ. Astron. Soc. Japan 19, 556.

Michaud, G.: 1970, Astrophys. J. 160, 641.

Michaud, G., Tarasick, D., Chariand, Y., Pelletier, C.: 1983, Astrophys. J.

269, 239.

Morgan W. W., Keenan, P. C., Kellman, E.: 1943, Atlas of Stellar Spectra,

University of Chicago Press, Chicago.

Parenago, P. P.: 1958, Soviet Astron. - A. J., 2, 151.

Sadakane, K., Takada, M., Jugaku, J.: 1983, Astrophys. J. 274, 261.

Sargent, W. L. W.: 1965, Astrophys. J. 142, 787.

Slettebak, A.: 1952, Astrophys. J. 115, 575.

Slettebak, A.: 1954, Astrophys. J. 119, 146.
-13-

Slettebak, A.: 1975, Astrophys. J. 197, 137.

Slettebak, A., Collins, G. W., It, Boyce, P B., White, N. M., Parkinson, T.

D.: 1975, Astrophys. J. Suppl. 29, 137.

Slettebak, A., Wright, R. R., Graham, J. A.: 1968, Astron. J. 73, 152.

co
+7

Table la. IUE observations of .a Boo and suspected a Boo-type stars

v sin i

Star HD Sp. Type (km/s) SWP LWP

,r l Ori 31295 AOp 110 21318 2098

2557

HR 4881 111786 A3p - 21945 2561

X Boo 125162 AOp 100 21319 2099

21320

e Sgr 169022 B9p 100 21316 2097

21317

29 Cyg 192640 Alp 90 21321 2100

15 And 221756 Alp 100 21944 2559

Note: Abt (1984), who first listed 15 And as a probably a Boo star, gives

A1Vp (Mg wk) as the spectral type with v sin i - 109 km/s. On a Perkins 72"

spectrogram of Oct. 23, 1985 I confirm the spectral type to be near Al, the

Mg II 4481 line does indeed appear weak for that type, and the rotational

velocity is in the neighborhood of 100 km/s.
(+TI

Table lb. IUE observations of standard stars

i

v sin i

Star HD Sp. Type (km/s) SWP LWR

6 Cas 8358 A5 III-IV 100 11214 9831

S Ari 11636 A5 V 70 - 8712

18 Tau 23324 B8 V 250 7923 6901

y Gem 47105 AO IV < 10 15542 12021

s Car 80007 Al IV 125 3383 2970

a UMa 89021 A2 IV 35 8187 7115

6 Leo 97633 A2 V 15 15505 11981

a Dra 123299 AO III 15 3840 3423

109 Vir 130109 AO V 330 - 1922

a Lyr 172167 AO V < 10 11213 9830

6 Cyg 186882 B9.5 III 140 14822 3041

• PsA 216956 A3 V 85 14417 11022

• Peg 218045 B9.5 III 120 2671 2386

i And 222173 B8 V 90 10376 9057
Table 2a. Spectrum lines measured in the IUE short-wavelength range

r

Element Multiplet Wavelength (^,)

Si III UV 4 1298.89, 1298.96

0 I UV 2 1302.17

C II UV 1 1334.53, 1335.71

N I UV 4 1492.63, 1492.82,

1494.68 t

Si II UV 2 1533.44

t

t

C I Uv 2 1656.27, 1656.93, 1657.01, 1657.38,

1657.91, 1658.12

Al II 2 1670.81

Ni II 5 1741.56

ti

N I UV 9 1742.72, 1742.73

1745.25, 1745.26

S I 2 1820.37

AA III 1 1862.78

Fe II 126 1.864.74

4
   7 , 

77t

yy

-,-

,

a

4

i'

Table 2b. Spectri:m lines measured in the IUE long -wavelength range

Element Multiplet Wavelength (^)

Mn II 1 2576.11

2605.70

Fe II 1 2585.88

2598.37 02599.39

Fe II 63 2739.54

Fe 11 62, 63 2749.18 (63), 2749.32 (62),

2749.48 (63)

Mg II 1 2795.52

2802.70

Mg II 2 2928.62

2936.50
.- t 7

CAPTION FOR FIGURES

FIGURES 1-24. - Equivalent widths W(R) of the ultraviolet lines versus

spectral types of the program and standard stars. The a Boo and suspecteda Boo stars are plotted as squares, with symbols designating each star as

follows: n - n l Ori, HR - HR 4881, a - ^ Boo, E - E Sgr, 29 - 29 Cyg,

15 - 15 And. A colon behind any symbol indicates that the equivalent width

measurement is particularly uncertain. The dashed line in each figure has

been drawn by eye through the standard-sta." points, as a guide to the general

trend of the line strengths for the standard stars; it has no physical meaning

beyond that.
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