Colliding Stellar Winds In the Eclipsing Wolf-Rayet Binary V444 Cygni

DOUGLAS N. BROWN1'

4

 and STEVEN N. SHORE

 2

-

3

'<

Received: ; Accepted:

1. Department of Astronomy, University of Washington

2. Astrophysics Research Center, Department of Physics, New Mexico Institute of

Mining and Technology

3. DEMIRM, Observatoire de Meudon

4. Guest Investigator, International Ultraviolet Explorer Satellite •

(BASfl-CB-162937) COLIIDING SIEXLAB HINDS IB

1EE £C1IFSING WCiF-BMET BJKAE * V4<i<4 CYGNI

(KasbingtoB Uriv.) 21 p CSCL 03B

ABSTRACT.

We have obtained high resolution spectra of V444 Cygni with the International Ultraviolet Explorer Satellite. These spectra span both eclipses and include one observation

at third quadrature. Together with seven archival spectra, they provide reasonably

complete phase coverage for the system. The variations in the P Cygni profiles of the

He II and N IV lines imply the existence of a low density region in the WR wind.

This region occupies a relatively narrow range of orbital phase coinciding with the

highest terminal velocities observed in C IV. We interpret these data as evidence of an

interaction region separating the winds of the O-star and Wolf-Rayet star.

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I. Introduction

The importance of the Wolf-Rayet binary V444 Cygni = HD 193576 (O6-8 111:4

WN4 AD; Walborn 1974) is that the system affords an almost unique laboratory for

studying the structure of a massive radiatively driven stellar wind. In consequence of

the atmospheric eclipse of the primary by the Wolf-Rayet (WR) secondary, the density

and velocity structures of the constituent stellar winds are encoded in the variation of

ultraviolet line profiles with orbital phase. V444 Cygni is the only eclipsing WR 4 O

binary which is bright enough to be conveniently studied at high resolution with the

International Ultraviolet Explorer (IUE).

Radial velocity studies (Wilson,1940,1942; Keeping,1947; Munch,1950) have established the orbital parameters: a = 40A0, A/oe — 25A/@, and MWN = 10M0. The

most recent photometric solution (Cherepashchuk and Khaliullin 1973; Cherepashchuk,

Eaton, and Khaliullin 1984, hereafter CEK), based on analysis of the eclipses at wavelengths between 2460 A and 3.5/j, yields the orbital inclination, t = 78°, radius of the

primary, flOe = 10AO; and the radius of the WR core, RWK = 2.9fiG. These values

are consistent with earlier photometric determinations (Kron and Gordon 1943,1950)

and with a polarimetric determination of the orbital inclination by Rudy and Kemp

(1978).

The unequal durations of primary and secondary eclipse have generally been

considered to imply that the WR core is surrounded by an extended, optically thin,

strongly limb-darkened outer atmosphere (Kopal, 1944) of radius roughly three times

that of the core (Russell, 1944). Kopal and Shapley (1946) predicted that electron

scattering was the principal source of continuous opacity in this dilute envelope, a fact

confirmed by the observation that the shape of primary eclipse is independent of the

wavelength of observation (Hiltner, 1949).

The mass loss rate has been determined by two methods. Bieging, et al. (1982)

and Abbott, et al. (1986) have measured the free-free emission at 4885 MHz. and

derived M = 1.4 x 10"sMovr"' on the basis of the theory of Wright and Barlow

(1975). Khaliullin (1974) and Kornilov and Cherepashchuk (1979) have used the rate

of change of the orbital period (see eq. 1) and the assumption of radial outflow from

the WR component to infer M = 1.2 x lO-'Mgyr"1.

CEK find that the color temperature of the WR core in the optical-ultraviolet

range is approximately 20,000 K, while the electron temperature is 80,000 - 100,000

K. The wide wavelength range of their data allowed CEK to determine the run of the

volume extinction coefficient with radial distance from the WR component. Attributing

the extinction to Thompson scattering and assuming that the electrons arise entirely

from the double ionization of helium, they determined that at the WR core, n. =

9 x 10'2cm~3 and v = 400 kms"1. They find that the wind velocity increases steeply,

attaining an outflow velocity near 2000 kms"1 at r = 10fla. The wind accelerates more

gradually at larger radii, attaining its terminal velocity (2500 kms*') at a distance of

20 to 30 R0.

The optical spectrum is dominated by high excitation emission lines of He II,

N III; N IV and N V which display a strikingly diverse range of phase dependent

variations (Munch, 1950). The blend of He II and N III at A4542 is stronger at primary

eclipse than at secondary and weakest at elongations, while N III A4640 is stronger

at secondary minimum than at primary. The flux in the N III A3483 emission feature

is roughly independent of orbital phase except at primary eclipse, when the emission

weakens markedly and a strong absorption feature appears. In contrast, N IV A4058

lacks absorption at any phase and is strongest at primary minimum.

Kuhi (1968a) obtained narrow-band light curves through secondary eclipse in
the principal optical emission lines. His results confirm the diversity of line variations

(some lines do not show eclipses at all) and also reveal the magnitude of intrinsic intensity fluctuations in the optical emission lines. Flux differences at constant phase

of order tens of percent were found, even when the observations were separated by

only a few orbital periods. Kuhi determined from his observations that the WR atmosphere, though optically thin in the continuum, is optically thick in the emission lines

and suggested that electron scattering of line radiation contributes significantly to the

(fluctuating) characteristics of the narrow band eclipse curves. He was accordingly

unable to account for his light curves on the basis of geometrical considerations alone.

From the general tendency for emission lines to be weaker at primary than at secondary

eclipse, he inferred that a significant fraction of the flux emitted in lines comes from a

small cylindrical region between the two stars.

Kuhi (1068b) and Bisiacchi, et al. (1982) found that the profile of He II A4686

is symmetrical at both elongations but asymmetrical to opposite sides during primary

and secondary eclipse; it is shaded to the blue at primary eclipse and shaded to the red

at secondary. Sahade (1958) reported the detection of a narrow emission component

superimposed on the broad A4686 emission feature. He interprets the radial velocity

variations of this component as indicative of an origin near the Lagrangian point, I,,

of the system. Underbill (1966) has suggested that the presence of sharp displaced

absorption features particularly near primary eclipse (9 - O.'OO) is evidence for gas

streaming in the system.

Recently, low dispersion IUE spectra have been employed to characterize the

variations of UV emission lines in this system. Eaton, Cherepashchuk and Khaliullin

(1985, hereafter ECK) noted that the C IV AA1548,1550 and Si IV AA1396,1403 resonance doublets showed P Cygni profiles with the largest terminal velocities and least

dependence on orbital phase, while the N IV A1718 subordinate line showed the lowest

terminal velocity and greatest phase dependence. Additionally, they noted that the He

II A1640 and N IV A1718 lines showed weak absorption components near secondary

eclipse. Synthetic photometry in these lines showed:

1. Primary flux minimum for He II, N IV, and Si IV are displaced from the

continuum mid-eclipse phase, occurring at * = 0.05.

2. In He II and N IV, the flux increases between phases 0.25 and 0.45.

3. The emission in all three of these lines decreases at secondary eclipse.

They interpret these light curves in terms of a model in which the emitting region

lies within 0 ± 4R0 of the WR component but the absorption arises in a spherical shell

of inner radius s= 10AO and outer radius external to the system. The brightening of

A1640 and A1718 between * = 0.25 and * = 0.45 (% 5 percent of the continuum flux) is

interpreted as a binary reflection effect: the O star radiation reflected off the expanding

WR envelope. In contrast, light curves in the N IV| A1486, C IV AA1548.50 and N V

A1240 display little if any phase dependence, whereby ECK infer that they are formed

exterior to the formation region of the He II and N IV lines.

Koenigsberger and Auer (1985) measured fluxes and equivalent widths in emission and absorption components of some P Cygni lines as well as velocities of identifiable

points on the line profiles. They determined that both the emission and.absorption

components of He II A1640 and N IV A1718 profiles vary with phase. Most striking was

their observation that the blue-shifted absorption of C IV A1550 is narrower at primary

eclipse than near secondary eclipse. They attributed this difference to the contribution

of an O-star wind which reaches a higher terminal velocity than that of the WR wind.

Accordingly, they suggested the likelihood of a wind-wind interaction region between

the two stars.
Favoring an alternative viewpoint, Underbill (1983a,b;1984 a,b;1986) has persistently criticised the widely accepted notion that WR stars possess relatively small hot

cores and powerful, radiatively driven winds. Underbill and Fahey (1987) argue that,

because of the presence of weak unresolved emission lines, the UV light curves derived

(as those of ECK) from low resolution IUE spectra do not represent true continuum

fluxes. They derive continuum fluxes from a series of 13 high-resolution IUE spectra of

V444 Cygni. Comparing these data with the optical light curves of CEK, they argue

for the presence of sufficient phase dependent "third light" shortward of 1500 A to invalidate the light curve analysis of Cherepashchuk and his collaborators. On the basis

of the orbital phase dependent variations of the He II A1640 profile, they recommend

a model in which the WR component has 7",// = 30,000 K, R = OfiO and possesses

large magnetically-supported plumes which contribute much of the optical and UV line

emission.

The original goal of the observations reported here was to determine the density

and velocity structure of the WR wind directly from the absorption components of the

ultraviolet P Cygni profiles. By subdividing the WR wind absorption profiles through

secondary eclipse into bins of ~ 300 km s~', we planned to generate a set of light curves,

each specific to a single projected expansion velocity. We would thus obtain the total

optical depth along iso-velocity surfaces for each impact parameter. In principle it is

possible to obtain, from the phase dependence of this optical depth, both the density

and velocity structure of the WR wind. In practice, the intended Doppler tomography

was incompletely realized because O-star absorption profile is not separable in the IUE

spectra, even at quadrature when the velocity difference between the components is

greatest. Nevertheless, the spectra obtained for this program reveal very important

features of the binary system.

II. Observations

The new observations discussed here were obtained using the S WP camera of the

IUE satellite during low noise (US1) shifts in 1085 May/June. The large aperture was

used for maximum throughput and because of the uncertain flux calibration of small

aperture high resolution spectra. The choice of exposure times was determined by the

need to achieve maximum phase resolution consistent with acceptable S/N at C IV

A1550. Because the star is relatively faint and heavily reddened, the longest exposure

which could be accommodated in a half-shift yielded S/N in the continuum at A1550

of roughly 5:1, only half that expected for an optimal exposure. These spectra were

supplemented by those archival images which were taken through the large aperture

and which have acceptable noise levels. Several archival images were rejected because

of poor S/N, although their omission left phase gaps. A journal of the new observations

is presented in Table 1, while corresponding data for the archival spectra appears in

Table 2.

The spectra are phased according to the ephemeris of Kornilov and Cherepashchuk

(1079)

HJD(primarycelipse) = 2441164.337 •+ 4.'212435(£) (1)

supplemented by a correction for the observed rate of period change (Cherepashchuk

1982). For SWP 14553, (O-C) = -O.'Oll, while (O - C) = -O.'OIS for all the other

images.

The data were analysed using routines available at the Goddard Space Flight

Center Regional Data Analysis Facility. Because the spectra are underexposed, smoothing beyond the optimal 3-point standard (which matches detector and spectrograph

resolution) was required. This is a delicate procedure, as excessive smoothing can not

only obscure important features but also introduce spurious ones. We performed a

8
number of comparison tests with a running boxcar filter, using the reseaux and the

profiles of the sharp interstellar C IV lines as fiducial indicators in this connection.

Easily perceptible changes in profile slope are produced by smoothing to more than 10

points; serious loss of information begins to occur with 20-point smoothing. We employ

either 5- or 10-point smoothing, depending upon the S/N of the original image. Those

archival spectra which were rejected because of inadequate S/N were found to require

smoothing to 20 or more points before they could be Intel-compared.

In the cases of the strongest lines, N V AA1239.43 and C IV AA1548,51, improper

interorder subtraction yields negative flux values at the bottoms of the saturated absorption troughs. The equivalent widths for these lines were computed by setting all

negative fluxes to zero. The unsaturated He II A1640 and N IV A1718 absorption lines

were free of this defect.

All the lines discussed below exhibit important periodic changes in line structure,

especially near the cross-over point between the emission and absorption components

of their P Cygni profiles. In many cases the line profile slope is greatest at precisely

this cross-over point. These facts make it difficult to define the appropriate continuum

level which separates the absorption and emission components, especially in relatively

noisy spectra. The local pseudo-continuum outside the line is much more easily defined,

whereby the wavelength interval for the total equivalent width (emission + absorption)

can be dependably selected. Accordingly we quote only latter quantity; any interpretation of these data which depends upon the detailed parsing of the P Cygni structure

is far less certain.

Finally, we infer the terminal velocity from the wavelength at which the extrapolated profile merges with the local pseudo-continuum. Careful tests of sample profiles

subjected to different degrees of smoothing confirmed that negligible error in the mea9

sured terminal velocity is introduced by the smoothing procedure unless the width of

the smoothing filter significantly exceeds that used in our reductions,

III. Results

Our analysis is based upon the variation of the four strongest P Cygni lines

with orbital phase. The A1640 and A1718 lines, because they arise from highly excited

states, are more sensitive indicators of wind structure within the binary orbit than are

the resonance lines. Themonotonic increase of SWP camera sensitivity with A > 1500A

multiplies this advantage. Conversely, the resonance lines (AA1239.43 and AA1548,51)

probe the velocity structure at large radii. The measured equivalent widths for the

principal lines discussed in this section are listed in Table 3.

a. He II A1640

The emission line always dominates this profile, whereby the total equivalent

width is negative. The absorption is strongest at primary eclipse, weak at secondary

eclipse and completely absent at phase 0.60. The single wave variation of Wwto with

orbital phase is shown in Figure 1.

Near primary eclipse the emission profile is rounded, symmetric and red-shifted

from the rest wavelength in the WR frame. The saturated absorption profile is also

rounded, mirroring the emission profile shape. Figure 2 shows the contrast between

the profile at primary eclipse and that at * = O.'OO.

Figure 3 shows gallery of line profiles which span the half of the orbit centered

on secondary eclipse. Near quadratures, the emission profile strengthens slightly, developing an overall triangular shape peaked at the WR velocity with a steepened blue

edge. The absorption profile also becomes triangular, giving a sawtooth appearance to

the P Cygni line. Near secondary eclipse, the emission profile becomes slightly broader,

10
distinctly plateau-shaped with the steep blue edge characteristic of the quadrature profiles, and symmetric with respect to the WR velocity. These variations of the emission

line profile do not extend to the high velocity red tail. Longward of approximately

A1645, the emission profile is invariant with orbital phase.

The terminal velocity inferred from the absorption line is clearly variable on time

scales ranging between hours and days. There is no evidence for coherent variation with

orbital phase, however. The terminal velocities, transformed to the binary center of

mass frame (7(06) = 10 km s~'), are shown in Table 4 and plotted against orbital

phase in Figure 4.

6. N IV A1718

A prominent absorption component, visible at most phases on the P Cygni line,

reaches its maximum at primary eclipse and is weakest near secondary eclipse. The

emission component is strongest near quadratures, when the absorption is also strong.

This alternation between absorption- and emission-dominated profiles contrasts sharply

with the character of the He II line, which is always strongly dominated by emission.

The total equivalent width of A1718, however, varies in phase with that of He II (see

Table 3 and Figure 5).

The N IV emission feature near primary eclipse is round-topped on some orbits,

flat-topped on others, but always centered significantly redward of the WR velocity.

When the emission is weakest, the absorption sometimes extends redward of the WR

velocity. The most important changes in the profile occur near secondary minimum,

when the absorption component weakens, completely disappearing at * = 0.P52. The

emission component also weakens beginning at * = 0."46, but recovers by * = 0.P76.

The occurrence of this almost pure emission profile is thus not phase-symmetric about

secondary eclipse.

11

Near first quadrature (SWP 24731; # = 0.'37), the emission and absorption

profiles are both triangular. The asymmetric emission is peaked at the WR velocity,

shaded to the red and with a steep blue edge. Near third quadrature, the emission

profile differs only in that it is shifted redward and truncated, forming a plateau-shaped

profile.

The terminal velocity (see Table 4 and Figure 4) is generally consistent with

that of He II, approximately constant on short timescales at about 2100 km s~'. The

absorption profile is generally asymmetric, displaying the steepest blue edge at the

quadratures. When strongest the absorption is nearly saturated. An extensive gallery

of A1718 line profiles was presented in a preliminary report of these observations (Brown

and Shore 1086). In Figure 6 we show the consistency of the profiles near primary

minimum and the contrasting secondary eclipse profile.

c. C/VAA1548.51

The total equivalent width of this line is almost always approximately zero (equal

fluxes in emission and absorption), suggestive of formation by scattering of light from

a small core embedded in an optically thick primarily scattering-medium. The deepest

absorption is seen during primary eclipse.

The profile variations are not striking at most phases, so that, as noted by

ECK, the low dispersion spectra supply little information. The emission profile varies

between two relatively stable shapes. Between phases 0.P37 and O.*60, the emission is

skewed slightly to the blue. A steep blue emission edge is joined to a correspondingly

steep red edge of the absoption profile. Between O.l>02 and O.'IO, the emission profile

is shallower and more symmetrical, blending into the absorption line with a relatively

shallow slope. Figure 7 displays the constancy of the profile through primary eclipse.

12
The transition between these forms occurs in a phase interval of approximately O.'l

near third quadrature; the phase gap at first quadrature prevents us from observing the

opposite transition. Except at those phases near secondary eclipse when the absorption

to emission transition is very steep, the absorption and emission components join at

1548 ± lA in the WR frame.

The equivalent width varies by ±4 A (see Figure 8). The mean terminal velocity

is 2620 ± 420 km s"1 (3.30 variation). The most dramatic changes occur in the blue

absorption wing near secondary eclipse. A high velocity blue-shaded tail develops at

this phase, reaching a terminal velocity of 3800 km s'1. The feature is present at phase

O.M4 and has disappeared by phase 0.'16. Figure 0 shows the phase variation of the

terminal velocity, while a gallery of line profiles spanning secondary eclipse is shown in

Figure 10.

The development of the high velocity tail is accompanied by a change in the

strength of the emission feature, which becomes slightly broader and shallower. While

this last effect may be real, it is more likely due to diminished low velocity absorption,

as the redward edge of the blue-shifted absorption component is steepest at this phase.

While the profile remains saturated at all phases, the total equivalent width of the line

becomes most strongly negative (emission dominating) at this phase rather than at

primary eclipse.

d. N VAA1239.43 A

Because this doublet is badly blended with Lyo and falls in a noisy portion of

the spectrum, the variations can be only roughly characterized. The structure of the

line is typical for the resonance line of a WR star. The profile is always saturated but

the structure of the absorption trough, unlike that of C IV, is entirely dominated by

13

noise and blends. The terminal velocity appears to be comparable to that exhibited

by the C IV line, and the same high velocity tail seems to develop at phase 0,P60. The

poor S/N of the data in this part of the camera, however, renders these judgments

uncertain.

III. Discussion

The C IV and N V resonance lines certainly contain contributions from both

components, but the He II and N IV lines appear in the spectra of O6 stars as unsaturated absorption lines with no hint of P Cygni emission (Walborn et al., 1986).

Therefore, all the emission and much of the absorption observed in these lines must

arise in the WR wind. We therefore discuss the interpretation of the resonance and

non-resonance lines separately.

a. He II and N IV

Assuming the CEK solution for the O-star and WR core radii, the well-determined

orbital inclination implies that only slightly more than half of the WR core is eclipsed

at secondary minimum. Accordingly, the more than two-fold weakening of the A1718

and A1640 absorption at that phase implies that either, (a) the CEK value of the WR

core radius is too large, so that the WR core is actually more substantially eclipsed

or, (b) these absorption lines are formed almost entirely by absorption and scattering

of the O-star (rather than the WR) continuum in the WR wind. We adopt the CEK

photometric solution and hypothesis (b), arriving at the following deductions.

1. The presence of the absorption features at quadrature implies that the WR

wind envelopes the entire system. ECK reached a similar conclusion on the basis of

their continuum light curves for A < 1500A.

14
2. The disappearance of the A1718 and A1640 absorption near secondary eclipse

implies a dramatically lower density of the WR wind in the line of sight at those phases.

ECK noted the weakening of the absorption features in their low dispersion spectra

and interpreted it as a result of a binary reflection effect. The reflection (presumably

resonant backscattering) of O-star continuum light by the (approaching) near side of the

WR wind was assumed responsible for filling in the absorption line and for augmenting

the emission profile, thus accounting for the observed hump in their narrow band light

curves. Our high resolution spectra are clearly inconsistent with this suggestion. It is

well known that the Doppler shift imposed by a "moving mirror" is twice the velocity

of the mirror. CEK determined that the wind emerges from the WR core at 500 km s"1

and reaches terminal velocity (=2000 km s'1 in the He II and N IV lines) at roughly

half the separation of the binary components. Resonant backscattering would thus

produce a broad enhancement of the continuum, extending between 1000 km s'1 and

4000 km s-' (4 - 20A) from the rest wavelength in the WR frame. The gallery of A1640

spectra in Figure 3 show no hint of such a feature. It is difficult to imagine a scattering

geometry which would so convincingly mimic a reduced column density in the line of

sight.

3. The fact that He II and N IV emission is strongest at quadrature implies

that the strongest emitting/scattering region is on the side of the WR wind facing the

O-star. The weakness of N IV emission near primary eclipse supports this inference.

4. The remarkable constancy of the He II emission profile through primary

eclipse confutes the suggestion of a material stream from the WR toward the O-star,

as such a stream would presumably produce a recognizable redshifted absorption at

these phases. The eight profiles which fall in the interval 0.P89 and 0.P09 are identical

within the photometric uncertainty. Redward of 500 km s"1 from the WR velocity

15

(roughly the velocity which CEK attribute to the WR wind at the photosphere) eve

the profiles at 0.P76 and O.'OO are identical.

Notwithstanding the uncertainty in identifying the cross-over point between the

emission and absorption components of the line profile, it is clear that the red absorption

edge is never less than 200 km s'1 from the WR velocity at any phase. The absence

of low velocity absorption suggests that the WR wind is optically thick and supersonic

even at the photosphere, as inferred by CEK from their IUE spectrophotometry.

The C IV profiles at third quadrature and at <t> = 0.'37 differ very little in

shape; they are almost identical except for a relative velocity shift of approximately

400 km s"1, somewhat less than the relative WR velocity difference between those two

phases. It appears that the contribution of the O-star to the P Cygni profile at these

phases must be slight, whereby the profile variations must be interpreted in terms of a

non-spherically symmetric wind structure.

The interval of orbital phase during which the high velocity tail develops in C IV

and the absorption almost disappears from the A1718 line is moderately well resolved

in our data. It is evidently asymmetric about secondary eclipse ( 0.P44 < * < O."60)

and appears to be sharply bounded in phase: one edge occurs between O.W and O.M4,

the other between O.»60 and 0.<>76. The absorption disappears from A1640 not at

secondary eclipse, but at » = O.'eO. These characteristics suggest a dynamical origin

for the phenomenology.

e. Wind - Wind Interaction

The simplest explanation is that we are seeing a cavity trailing behind the O

star: a shadow in the WR wind, produced by the interaction of the winds of the two

16
stars. If Walborn's (1074) luminosity classification (O8-8 III) is correct, the O-star

likely possesses a more substantial wind than typical for a main sequence star of its

mass. A standing shock would accordingly be expected at the interface between the

two winds. We interpret the high velocity tail of the C IV resonance line as evidence of

an oblique shock at the boundary of the interaction region and the absence of the He

II and N IV absorption as evidence that the O star effectively shadows the WR wind.

When the cavity crosses our line of sight, we see the lower density O-star wind. In

the resonance lines, we see the asymptotic material velocity in the interaction region

projected onto the line of sight. The emission is, as usual, interpreted as the result of

scattering the O-star light off the WR wind.

The relative orbital velocity of the O-star with respect to the WR wind yields

an aberration angle of approximately O.tM between the line of centers and the axis

of the "wind shadow" cavity. This value is consistent with the observed phase offset

of the spectral indicators with respect to secondary eclipse. The opening angle of the

wind cavity appears to be approximately 60° so that the line of sight Is parallel to the

asymptotic conical surface of the interaction region at only two phases: shortly before

and after secondary eclipse. The resonance lines should exhibit the highest velocity

blue-shifted absorption at these phases. In the interval 0.P45 - 0.P60, the "terminal"

velocity observed in these lines will be reduced by cos\(W - i) - 30°] a 0.9. Figures

11 and 12 show a pair of cartoons which represent the geometry suggested by these

considerations. Further details of the model are discussed in the next section.

Underbill and Fahey (1087, hereafter UF) have proposed a very different model

for V444 Cygni on the basis of a diverse set of high resolution IUE spectra, many

obtained during the higher noise US2 shift. An important element of their critique

of the CEK photometric solution was noted above: the imputed presence of a phase17

dependent "third light" in the system. The ultraviolet continuum fluxes which UF

derive from high resolution spectra are certainly less affected by unresolved weak emission lines than are the CEK fluxes. The high resolution calibration, however is far

less secure than that of the low resolution spectra used by CEK. The high resolution

photometric calibration is poorest at the short wavelength end of the SWP camera,

where crowding of the echelle orders complicates the background subtraction. The

fluxes derived from even optimally exposed spectra are dependable to no better than

10%. For non-optimal images, other difficulties diminish the photometric accuracy.

Oliversen (1084) found that near 1300A, where UF find the greatest divergence from

the CEK light curve, linearity errors of 10 - 20% are not uncommon in trailed low

resolution spectra if, like the spectra of UF, they are underexposed and degraded by

high particle background. High resolution spectra may be expected to suffer (at least)

comparable linearity errors.

We have supplemented our US1 observations with those of the UF spectra which

have roughly comparable S/N. Yet, we find no evidence in the resulting set of A1640

or A1718 profiles for the compact plasma tubes which they suggest link the two stars.

For example, the peaks of the A1640 profiles in Figure 3 track the WR star velocity.

One possible origin of the features which they discuss is the smoothing procedure

which they employ. It is important to note that a 51-point smoothing is not equivalent

to binning with IOA resolution because the smoothing algorithm performs a running

average centered on each pixel. In order to improve S/N at the expense of resolution,

it is necessary to average the data in non-overlapping wavelength bins. This procedure

avoids the generation of correlated errors in independent pixels.

IV. Toward a Model for V444 Cygnl

Interacting winds have been received relatively little attention in the literature.

18
Cherepashchuk (1076) and Prilutskii and Usov (1076) have presented qualitative discussions applicable to WR binaries. Chapman (1981), Huang and Weigert (1082) and

Girard and Willson (1087) have undertaken more extensive analysis, but for systems

which bear little resemblence to V444 Cygni.

For simplicity, we assume that the winds of both stars have reached terminal

velocity prior to collision. The most recent models of Pauldrach et al. (1086) suggest

that this is the case for the WR wind in V444 Cygni, but it may be only approximately

true for the O-star wind. The location of the stagnation point is determined by the

equality of the wind dynamical pressures. Taking o = (MjUj/Mifi)1'

2

 gives x = (jf;)

for the position of the stagnation point with respect to the O-star in terms of the

binary separation. Since the velocities of both winds fvery high, we neglect the effects

of orbital motion except for a slight aberration from the line of centers. There is no

evidence that the WR wind is directly impinging on the surface of the O star, so

that Q must be sufficiently small that the stagnation point is well separated from its

photosphere. The opening (full cone) angle of the shock relative to the WR star is

about 6 — 40° so that an approximate solution gives z =s 0.3.

This places the stagnation point close to the center of mass of the system and

serves as a constraint on the mass loss rate from the O star. The derived value is

a = 0.3 — 0.5. Assuming that the wind velocities are approximately equal, MO ~

(1 — 3) x 10~6M0j/r~'. This value is far closer to that of an evolved than main sequence

O star, consistent with the spectroscopic classification of Walborn (1074). We also note

that none of the ultraviolet spectral features conflicts with Walborn's luminosity class

assignment.

IV. Conclusion

We interpret the abrupt increase near secondary eclipse in the A1550 terminal

10

velocity and simultaneous decrease in the material density indicated by A1640 and

A1718 absorption as evidence for a shock dominated wind-wind interaction between

the binary components of V444 Cygni. Our high resolution IUE spectra provide no

evidence for the existence of material streaming between the two stars, as suggested by

Kuhi (1068a) or of magnetically configured plasma plumes and tubes, as suggested by

Underbill and Fahey (1087).

Because the momenta of the two winds are large and the hydrodynamics of

their collision is complex, only a preliminary outline of a model is currently possible.

A qualitative prediction is, however, implied. Of the four other currently known (WR

-I- O) eclipsing binaries, only CQ Cep should exhibit similar phenomenology in IUE

spectra. The separation of the components of CV Ser is likely too large in comparison

with the wind momenta for a similar structure to be produced. GP Cep is a relatively

compact quadruple system in which the winds of all four components may play a role.

CX Cep is much too faint for IUE observations. Stickland et al. (1084) noted an

apparent depletion of He II and N IV at primary eclipse, when the leeward side of the

presumed O-star companion crosses the line of sight. We predict that high resolution

observations will reveal an abrupt increase in the A1550 terminal velocity at phases

near the eclipse.

A second implication of the proposed model is that detectable linear polarization

would be expected in the UV emission lines because of the presumed importance of

resonance scattering in their formation. The optical thickness of the lines contrasts with

the optically thin continuum, rendering a quantitative prediction beyond the scope of

this paper.

The importance of V444 Cygni as a unique astronomical laboratory for the study

of steady state thermodynamics and hydrodynamics in colliding stellar winds amply

20
justifies the theoretical and observational effort which will be required to refine the

model and test its predictions.

Acknowledgements

We wish to thank Drs. N. Bochkarev, C. Catala, J. Castor, J. Eaton, R. Fahey,

R. Hjellming, G. Koenigsberger, Y. Kondo, J. Sahade, G. Sonneborn, A. Torres.and

N. Walborn for enjoyable discussions about Wolf-Rayet systems in general and our

study in particular. SNS warmly thanks Dr. P. Encrenaz and Dr. F. Praderie for

their invitation to the Observatoire de Meudon during Summer 1986 when some of the

preliminary work in this study was carried out. This work is supported by NASA grant

NAG 5-357 to the University of Washington.

21

REFERENCES

Abbott, D.C., Sieging, J.H., Churchwell.E., Torres, A. 1986, Ap. J., SOS, 239.

Bieging, J.H., Abbott, D.C., and Churchwell, E.B. 1982, Ap.J., 263, 207.

Bisiacchi, G.F., Firmani, C., and de Lara, E., 1982, in Wolf- Rayet Start: Observations,

Physics and Evolution, IAU Symposium No. 99, ed. de Loore, C. and Willis, A.

(Dordrecht; Reidel), p. 283

. Brown, D. and Shore, S. 1986, in New Insights in Astrophysics, Proc. Joint NASA/ESA/SERC

Conf., ESA SP-263, p. 353.

Castor, J., Abbott, D. and Klein, R. 1975, Ap.J., 195, 157.

Castor, J. and Lamers, H. 1979, Ap.J.Suppl., 39, 481 .

Castor, J., Lutz, J. and Seaton, M. 1981, MNRAS, 194, 547.

Chapman, C.R. 1981, Ap.J., 248, 1043

Cherepashachuk, A.M. 1976, Sov.Astr.Lctt., 2, 138.

1982, Aitr.Sp.Sci., 86, 299.

Cherepashchuk, A. and Khaliullin, Kh. 1973, Sov.Astr.-AJ, 17, 330.

Cherepashchuk, A., Eaton, J., and Khaliullin, Kh. 1984, Ap.J., 281, 774 (CEK).

Eaton, J., Cherepashchuk, A. and Khaliullin, Kh. 1985, Ap.J., 297, 266 (ECK).

Girard, T. and Willson, L. 1987, preprint.

Hamman, W-R. 1985, A air. Ap., 145, 443.

Hiltner, W.A. 1949, Ap.J., 110, 95.

Huang, R. and Weigert, A. 1982, Astr. Ap., 112, 281.

Keeping, E.S. 1947, Publ. Dom. Ap. Obi., 7, 349.

Khalliullin, Kh. 1974, Sov. Astr.-AJ, 18, 229.

Khalliullin, Kh. and Cherepashchuk, A. 1976, Sov. Astr.-AJ, 20, 186.

Koenigsberger, G. 1984, Rev.Ap.Mei., 9, 159.

22
Koenigsberger, G. and Auer, L. 1985, Ap.J., 297, 255.

Kopal, Z. 1944, Ap.J., 100, 204.

Kopal, Z. and Shapley, M. 1946, Ap.J., 104, 160.

Kornilov, V. and Cherepashchuk, A. 1979, Sov. Astr. Lett., 6, 214.

Kron, G. and Gordon, K. 1943, Ap.J., 97, 311.

— 1950, Ap.J., Ill, 454.

Kuhi, L. 1968a, Ap.J., 152, 89.

1968b, in Wolf-Rayet Start, ed. Gebbie, K. and Thomas, R., Nat. Bur. Stds.

Spec. Publ. 307, (Washington, D.C.: U.S. Gov. Print. Off.)

Munch, G. 1950, Ap.J., 112, 266.

Oliversen, N. 1984, NASA 1UE Newsletter No. 24, p. 27.

Pauldrach, A., Puls, 3., Hummer, D., and Kudritzki, R., 1985, Astr. Ap., 148, LI.

Prilutskii, O. and Usov, V. 1976, Sov. Aitr.-AJ, 20, 2.

Rudy, R. and Kemp, J. 1978, Ap.J., 221, 200.

Russell, H.N. 1944, Ap.J., 100, 213.

Sahade, J. 1958. Mem. Soc. Roy. Sei. Liege, 20, 46.

Stickland, D., Bromage, G., Budding, E., Burton, W., Howarth, I., Jameson, R., Sherrington, M., and Willis, A., 1984, Astr. Ap., 134, 45.

Underbill, A. 1966, The Early Type Stan (Dordrecht:Reidel)

1983a, Ap.J., 265, 933.

1983b, Ap.J., 266, 718.

1984a, Ap.J., 276, 583.

1984b, Ap.J., 287, 874.

Underhill, A. and Fahey, R. 1987, Ap.J., 313, 358 (UF).

Walborn, N. 1974, Ap.J., 189, 269.

23

Walborn, N., Nichols-Bohlin, J. and Panek, R. 1985, International Ultraviolet Explorer

Atlai of O-Type Spectra, NASA Ref. Publ. 1155

Wilson, O.C. 1940, Ap.J., 91, 379.

1942, Ap.J., 95, 402.

Wright, A. and Barlow, M. 1975, M.N.R.A.S., 170, 41.

24
Table 1. Journal of New Observations

SWP

25962

25990

25994

26000

26007

26008

26030

26031

26041

26065

26066

Exposure (minutes)

192

245

200

192

200

205

200

205

202

200

200

JD - 2,440,000

6202.884

6207.023

6208.889

6209.891

6210.886

6211.248

6214.890

6215.051

6217.050

6218.812

6218.968

*

0.098

0.080

0.523

0.761

0.997

0.036

0.948

0.986

0.461

0.870

0.916

Table 2. Journal of Archival Spectra

SWP

14553°

24731'

24734'

24740'

24745'

25904'

25905'

Exposure (minutes)

90

170

180

105

200

200

205

JD - 2,440,000

4811.396

6056.580

6057.578

6058.587

6059.583

6195.887

6196.049

*

0.773

0.366

0.603

0.843

0.079

0.437

0.475

Notes:

' VILSPA image.

' US2 image.

' US1 image.

25

26
Table 3. Equivalent Widths (A)

SWP

14553

24731

24734

24740

24745

25904

25905

25962

25990

25994

26000

26007

26008

26030

26031

26041

26065

26066

*

0.773

0.366

0.603

0.843

0.079

0.437

0.475

0.098

0.080

0.523

0.761

0.997

0.036

0.948

0.986

0.461

0.879

0.916

N V 1240

2.8

0.48

-3.0

2.7

3.0

1.2

4.2

1.3

2.8

5.3

4.3

-3.7

-12.2:

-21.1:

-6.8

1.7

-3.7

-6.5

C IV 1550

-2.1

-0.8

2.8

-1.8

0.6

-0.5

1.1

1.8

3.1

3.5

-3.3

0.7

-1.0

2.6

0.6

-1.1

-0.0

0.6

He II 1640

-23.9

-18.9

-25.4

-20.2

-15.7

-19.2

-17.1

-16.1

-15.3

-20.7

-21.3

-19.3

-14.3

-12.8

-15.6

-21.1

-13.9

-13.1

N IV 1718

-2.4

-3.7

-5.8

0.8

3.1

-3.3

-3.5

2.1

3.3

-4.4

-3.2

1.7

3.0

4.7

1.8

-2.2

1.0

4.0

27

Table 4. Terminal Velocities (km s'1)

SWP

14553

24731

24734

24740

24745

25904

25905

25962

25990

25994

26000

26007

26008

26030

26031

26041

26065

26066

»

0.773

0.366

0.603

0.843

0.079

0.437

0.475

0.098

0.080

0.523

0.761

0.997

0.036

0.948

0.986

0.461

0.879

0.916

v,(06)°

110

-90

72

100

-58

-46

-19

-69

-58

17

120 (ff£

2

-27

38

10

-29

83

60

v,(WR)'

-292

238

-174

-244

155

128

58

186

157

-34

_^(j!67^

4

78

-88

-17

84

-200

-144

u,(1550)

-2380

-2340

-3480

-2270

-2100

-3230

-3050

-2180

-2180

-3370

^,Cj820;

-2410

-2310

-3290

-2290

-3810

-2310

-2260

v,(1640)

-2040

-2050

-1170

-1890

-1750

-1780

-2330

-2400

-2170

-1980

_j,(\j&t)

v,(1718)

-2300

-1590

-1810

-2570

-2520

-1930

-1640

-2100

-1820

-2180

jt

-1910 | -1910

-1880

-2760

-2440

-1980

-2440

-2340

-1980

-2560

-2530

-2030

-2480

-2250

Note:

° Derived from Munch (1950).

28
DOUGLAS N. BROWN:'Department of Astronomy FM-20, University of Washington,

Seattle, WA 08195

STEVEN N. SHORE: Astrophysics Research Center, Department of Physics, New Mexico Institute ol Mining and Technology, Socorro, NM 87801

29

FIGURE CAPTIONS

figure 1. Orbital phase variation of He II A1640 total equivalent width.

The uncertainty In phase Is less than the size of the plotting

symbols.

Figure 2. Comparison of four He II X1640 line profiles near primary

eclipse with that at 0?60 (SUP 24734). The eclipse profiles

are from SUP 26007, 26006, 26030 and 26031. Note the complete

absence of absorption at Of 60.

Figure 3. Gallery of He II X\6AQ line profiles spanning secondary eclipse.

Phase advances from the bottom (0?37) to the top (Of84) of the

figure. Note that the absorption component Is almost absent between

0.44 and 0?76.

Figure 4. Orbital phase variation of terminal velocities measured In

He II >1640 (t-rLa.*]/<4 and N IV Xl718 (dots). The uncertainty

In phase is less than the size of the plotting symbols.

Figure 5. Orbital phase variation of N IV 7«1718 total equivalent width.

Figure 6. (a) Overplot of N IV A1718 line profiles near primary eclipse.

The spectra are taken from the same Images as in Figure 2.

(b) Comparison of N IV A1718 line profile at secondary eclipse

(SUP 25994, 0.52) and just before onset of primary eclipse

(SWP 26065, 0?88). Compare with Figure 3) note that the

absorption component of >1718 disappears at earlier phase

than does that of %1640.

.Figure 7. Overplot of C IV A1550 line profiles near primary eclipse.

The spectra are taken from the same Images as in Figure 2.

Figure 8. Orbital phase variation of C IV Al550 total equivalent width.

Figure 9. Orbital phase variation of C IV X1550 terminal velocity. Note

that the Increased wind velocity coincides in phase with the

weakening of He II and N IV absorption.

Figure 10. Gallery of C IVX1550 line profiles spanning secondary

eclipse. Phase advances from the bottom (Of37) to the top

(0?84) of the figure. Note the presence of enhanced

blueshlfted absoptlon between oC*4 (SWP 25904) and of&0 (SWP 26000).

Figure 11. Cartoon of the proposed wind - wind interaction in V444 Cygnl.

Top view along the normal to the orbital plane. The SWP image

numbers of spectra discussed in the text are Indicated at the

corresponding orbital phases on the reference circle.

Figure 12. Cartoon side views of the V444 Cygnl system at secondary eclipse

showing the proposed structure of the wind - wind Interaction.

Direction of observer Is to the right.

(a) View along the line of nodes.

(b) View along the line of sight.
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