Radio emission from WR140

Abstract

Milliarcsecond resolution Very Long Baseline Array (VLBA) observations of the archetype WR+O star colliding-wind binary (CWB) system WR140 have been obtained at 23 epochs between orbital phases 0.74 to 0.97. The emission in the wind-collision region (WCR) is resolved as a bow-shaped arc which rotates as the orbital phase progresses. This rotation provides for the first time the inclination of the orbit (122+/-5 degrees), the longitude of the ascending node (353+/-3 degrees), and the orbit semi-major axis (9.0+/-0.5 mas). The implied distance is 1.85+/-0.16 kpc, which requires the O star to be a supergiant, and leads to a wind-momentum ratio of 0.22. Quasi-simultaneous Very Large Array (VLA) observations show the synchrotron spectra evolve dramatically through the orbital phases observed, exhibiting both optically thin and optically thick emission. The optically thin emission maintains a spectral index of -0.5, as expected from diffusive shock acceleration.

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arXiv:astro-ph/0410211v1 7 Oct 2004
**FULL TITLE**
ASP Conference Series, Vol. **VOLUME**, **YEAR OF PUBLICATION**
**NAMES OF EDITORS**
Radio emission from WR140
S.M. Dougherty
National Research Council, D.R.A.O, P.O. Box 248, Penticton, BC, Canada
A.J. Beasley
NRAO-AUI (ALMA), Camino del Observatorio 1515,
Las Condes, Santiago, Chile
M.J. Claussen
NRAO, 1003 Lopezville Rd., Socorro NM 87801, USA
B.A. Zauderer
Deptartment of Astronomy, University of Maryland,
College Park, MD 20742, USA
N.J. Bolingbroke
Department of Physics and Astronomy, University of Victoria,
3800 Finnerty Rd, Victoria, BC, Canada
Abstract.
tions of the archetype WR+O star colliding-wind binary (CWB) system WR140 have
been obtained at 23 epochs between orbital phases 0.74 to 0.97.
the wind-collision region (WCR) is resolved as a bow-shaped arc which rotates as the
orbital phase progresses. This rotation provides for the first time the inclination of
the orbit (122◦± 5◦), the longitude of the ascending node (353◦± 3◦), and the orbit
semi-major axis (9.0 ± 0.5 mas). The implied distance is 1.85 ± 0.16 kpc, which re-
quires the O star to be a supergiant, and leads to a wind-momentum ratio of 0.22.
Quasi-simultaneous Very Large Array (VLA) observations show the synchrotron spec-
tra evolve dramatically through the orbital phases observed, exhibiting both optically
thin and optically thick emission. The optically thin emission maintains a spectral
index of −0.5, as expected from diffusive shock acceleration.
Milliarcsecond resolution Very Long Baseline Array (VLBA) observa-
The emission in
The archetype of CWB systems is the 7.9-year period WR+O system WR140
(HD193793). Its highly eccentric orbit (e ≈ 0.88) modulates the dramatic variations
in the emission from the system observed at many wavelengths. At radio wavelengths
there is a slow rise from a low thermal state close to periastron of a few mJy, to a
frequency-dependent peak in emission of 10’s of mJy between orbital phase 0.65 to
0.85, before a precipitous decline just before periastron. The radio variations have been
widely attributed to an underlying synchrotron source viewed through the changing
free-free opacity of the extended stellar winds of the binary system along the line-of-
sight to the WCR as the orbit progresses (Williams et al. 1990; White & Becker 1995).
However, none of the free-free opacity models explain the radio light curve in a sat-
isfactory manner. Models that include processes intrinsic to the WCR are now being
explored (Pittard et al., in preparation).
We report briefly on high resolution observations of WR140 obtained with the
VLBA that image structures in the WCR at a linear resolution of a few AU, ap-
proximately the stellar separation at periastron, along with quasi-simultaneous VLA
observations. Dougherty et al. (2004) include a more detailed description of this work.
1

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Dougherty, Beasley, Claussen, Zauderer, Bolingbroke
Figure 1.
The contour levels are -1, 1, 1.6, 2.6, 4.1, 6.6, 10.5ρ where ρ = 220 µJy beam−1. The contour
levels and greyscale are identical in each image. The beam size is shown in the lower left
corner of the images, and is typically 2.0 × 1.5 mas2, which at a distance of 1.5 kpc gives a
linear resolution of 3 AU. Rotation and proper motion of the WCR are clear.
8.4-GHz VLBA observations of WR140 at orbital phases 0.74, 0.86 and 0.93.
1.Observations
Observations of WR140 were obtained at 23 epochs using the VLBA. The campaign
started on Jan 4, 1999 near the peak of radio emission around orbital phase 0.75 and
was completed Nov 18, 2000 when the radio emission had declined to its low level. The
resulting phase-referenced 8.4-GHz images at three of the observed epochs are shown in
Fig. 1. Additionally, we also obtained closely concurrent observations with the VLA at
five frequencies between 1.4 and 22 GHz.
The 8.4-GHz emission detected by the VLBA is clearly resolved. We identify this
emission as arising from the WCR in WR140 since this is the only emission in the
system with sufficient brightness temperature (> 105K) to be detected by the VLBA.
The stellar winds, with brightness of ∼ 104K are undetected. A bow-shaped ridge of
emission is observed at most epochs, as anticipated for the WCR from model calculations
(e.g. see Eichler & Usov 1993; Canto et al. 1996; Dougherty et al. 2003), with the bow
shock wrapping around the star with the lower wind momentum - typically the O star.
Between orbit phase 0.74 and 0.95, the WCR exhibits rotation from “pointing” NW to
W, in addition to an east to west proper motion of ∼ 10 mas.
2.Orbital parameters of WR140
Many of the orbital parameters in WR140 are well-determined from radial velocity
measurements (see Marchenko et al. 2003, and references therein). However, the orbital
inclination (i), semi-major axis (a) and the longitude of the ascending node (Ω) can
only be determined if the system can be resolved into a “visual” binary around the
orbit. The two stellar components in WR140 have been resolved using the Infrared-
Optical Telescope Array (IOTA) interferometer on June 17, 2003 to have a separation of
12.9+0.5
−1.3degrees east of north (Monnier et al. 2004).
Assuming P = 2899 days, To= 2446147.4, e = 0.881 and ω = 47◦(Marchenko et al.
2003), this single epoch observation at orbital phase 0.297 gives families of possible
solutions for (i,Ω,a).
Currently, the VLBA observations of the WCR are the only means to determine
uniquely i, and hence Ω and a, from the possible IOTA solutions. Under the assumption
that the free-free opacity along the line-of-sight to the WCR is sufficiently low as to not
impact the apparent distribution of emission from the WCR, we expect the “arc” of
WCR emission to wrap around the star with the lower wind momentum - the O star.
In this case, the rotation of the WCR as the orbit progresses implies that O star moves
from the SE to close to due E of the WR star over the period of the VLBA observations.
Also, if it is assumed the axis of symmetry of the emission from the WCR is coincident
−0.4mas at a position angle of 151.7+1.8

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Radio emission from WR140
3
Figure 2.
ω = 47◦, Ω = 353◦, i = 122◦and a = 9.0 mas at orbital phase 0.737, 0.858 and
0.931, overlaid on the VLBA 8.4 GHz images. The WR star is to the W (right)
of the WCR at these phases. The rotation of the WCR as the orbit progresses is
clear. The relative position of the stars to the WCR was determined using a wind
momentum ratio of 0.22 (see § 3.)
The derived orbit of WR140 on the plane of the sky using e = 0.88,
with the projection on the plane of the sky of the line-of-centres of the two stars, we can
derive the orbital inclination from the change in the orientation of the WCR with orbital
phase. Each (i,Ω) family provides a unique set of position angles for the projected line-
of-centres as a function of orbital phase. By fitting the position angle of the line of
symmetry of the WCR as a function of orbit phase for different sets of (i,Ω), we find a
best-fit solution of i = 122◦±5◦and Ω = 353◦±3◦. These values lead to a semi-major
axis of a = 9.0 ± 0.5 mas, and a projected orbit that evolves as shown in Fig. 2.
The derived orbit inclination is consistent with values previously suggested. How-
ever, it presents a challenge to current models of dust formation. To date, most models
of dust formation in WR140 assume that the gas from which dust is formed in the
WCR is compressed within ∼ 0.15 yr of periastron passage (Williams et al. 1990). The
subsequent motion of the compressed gas is determined by the velocity of this material
when it is compressed. Since the momentum of the WR star wind is higher than that of
the O star, this material moves away from the WR star, along the WCR. With the orbit
orientation derived here, the O star is NW of the WR star during periastron, and mate-
rial compressed at periastron will therefore have a proper motion to the NW. However,
high-resolution IR observations show that dust ejected during the 2001 periastron pas-
sage has proper motion to the south and east, away from the WR star (Monnier et al.
2002). New dust models are now attempting to address this challenge (Williams, this
volume).
3.Basic system parameters of WR140
Distance estimates of WR stars are typically based on absolute magnitude calibrations
that often have large scatter. Having determined the orbital inclination and semi-major
axis it is now possible to make an estimate of the distance to WR140 independent of
any stellar parameters. Marchenko et al. (2003) determined a sini = 14.10 ± 0.54 AU,
which leads to a = 16.6±1.1 AU for i = 122◦±5◦. Along with the derived semi-major
axis of a = 9.0 ± 0.5 mas, these give a distance of 1.85 ± 0.16 kpc.
This distance is somewhat larger than the usually quoted value of 1.3 kpc deduced
by Williams et al. (1990) from the luminosity of the system. Since the primary O-star
luminosity indicator is masked by the WC7 spectrum, Williams et al. (1990) assumed a
main sequence O4-5 star with an absolute magnitude of −5.6 and took that of the WC7
star to be −4.8. With the system at 1.85 kpc, the absolute magnitude of the O4-5 star
becomes −6.4, suggesting it is a supergiant (see Vacca et al. 1996, Table 7).

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Dougherty, Beasley, Claussen, Zauderer, Bolingbroke
Table 1.
Parameter
Distance (kpc)
Mv
Spectral type
BC
Mbol
log (Lbol/L⊙)
Mass (M⊙)
v∞(km s−1)
˙M (M⊙yr−1)
aBased on Mv and the calibration of Vacca et al. (1996)
M⊙,bol= 4.75m(Allen 1976)dSetia Gunawan et al. (2001)eEenens & Williams (1994)
Basic parameters of WR140
PrimarySecondary System
1.85
-6.8-6.4
0.4-5 Ia
-4.3a
-10.7
6.18c
54 ± 10
3100d
8.7 × 10−6
-5.6
WC7
-3.4b
-9.0
5.5c
20 ± 4
2860e
4.3 × 10−5
bWilliams et al. (1990)
cCalculated assuming
With the increase in distance, a reassessment of the mass-loss rates of the two stars
is appropriate. Based on the X-ray luminosity measured by ASCA (Zhekov & Skinner
2000), the mass-loss rate for the WC star at 1.85 kpc is 4.3×10−5M⊙yr−1. Repolust et al.
(2004) suggests values of 8.6 − 8.8 × 10−6M⊙ yr−1for O4-5 supergiants.
with the wind speeds (Table 1) these mass-loss rates imply a wind momentum ratio
η = 0.22. This wind-momentum ratio is considerably higher than the 0.035 deduced
by Williams et al. (1990). The higher value of η derived here, however, implies a half-
opening angle of the WCR of 63◦(following Eichler & Usov (1993)), consistent with
65◦± 10◦derived from these VLBA observations.
Along
4.The radio spectra of WR140
The new VLA data at five frequencies allow us to observe the radio spectrum and its
evolution better than previously possible, most particularly at the higher frequencies.
At phase 0.974, the spectrum is a power-law with a spectral index of 0.72±0.03, a value
characteristic of the stellar winds in WR+OB binary systems. Furthermore, the flux
levels at this phase are consistent with the wind densities implied by the parameters in
Table 1. Assuming the thermal emission from WR140 is essentially constant throughout
the orbit, the synchrotron spectra at each observed phase can be determined by simply
subtracting the thermal flux at phase 0.974 from the total flux. The resulting spectra
are shown in Fig. 3.
The synchrotron spectra between phases 0.67 and 0.92 are optically thin at several
frequencies, with a spectral index that appears to be closely constant, with a slope
of −0.5 ± 0.1, as expected for diffusive shock acceleration of electrons in strong, non-
relativistic shocks (e.g. Bell 1978; Drury 1983; Jones & Ellison 1991, and references
therein). The optically thick spectrum apparent during the bulk of the orbit has been
widely attributed to free-free absorption in the stellar winds along the line-of-sight to
the WCR. Unfortunately, these models are too simple to explain the radio observations
of WR140, as readily acknowledged by their authors. The VLBA observations (Fig. 1)
show the WCR as a distributed emission region and the lines-of-sight to the WCR
traverse different regions of the stellar winds. As a result, the emerging emission will
be a combination of both optically thick and thin emission since even though lines-of-
sight to the apex may be optically thick, a substantial amount of emission arises from
optically thin lines-of-sight to the downstream flow. Using the newly derived orbit and
assuming a half-opening angle for the WCR of 63◦, we now know the lines-of-sight to
the WCR traverse the O-star wind between orbit phases 0.24 and 0.99 during which
the most dramatic changes in the radio emission are observed. Clearly, if stellar wind
free-free opacity plays any role in determining the observed spectra, it is the O-star
wind opacity that is of most concern, not that of the WR star. Another shortcoming

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Radio emission from WR140
5
1
10
1 10
Flux (mJy)
Frequency (GHz)
0.671
0.800
0.831
0.848
0.907
0.921
0.930
0.955
Figure 3.
subtracting the thermal spectrum at phase 0.974 from the spectra at each phase.
The key to the plot is shown in the upper right-hand corner.
Synchrotron spectra of WR140 at several orbital phases, determined by
of previous models is the assumption that synchrotron emission is optically thin at all
observing frequencies, and at all orbital phases. The optically thick component of the
spectra may, at least in part, be due to mechanisms intrinsic to the WCR.
New radiative transfer models, based on a fully consistent hydrodynamic treatment
of the WCR, have started to explore the impact of a number of processes on the radio
emission from CWBs, including free-free opacity in the stellar winds and the WCR,
synchrotron self-absorption, Coulombic cooling through interactions with post-shock
ions, plasma effects such as the Razin effect, and Inverse Compton cooling by the in-
tense ultra-violet radiation field of the nearby massive stars. These models have been
very successful in explaining the radio emission from very wide CWBs such as WR147
(Dougherty et al. 2003), and are now maturing to the point where they will provide
more insight to the mechanisms acting in WR140 (Pittard et al., in preparation). The
observations presented here represent the essential constraints for these new models.
References
Allen, C. W. 1976, Astrophysical Quantities (London: Athlone (3rd edition), 1976)
Bell, A. R. 1978, MNRAS, 182, 147
Canto, J., Raga, A. C., & Wilkin, F. P. 1996, ApJ, 469, 729
Dougherty, S. M., Beasley, A. J., Claussen, M. J., Zauderer, B. A., & Bolingbroke, N. J. 2004,
ApJ, submitted
Dougherty, S. M., Pittard, J. M., Kasian, L., Coker, R. F., Williams, P. M., & Lloyd, H. M.
2003, A&A, 409, 217
Drury, L. O. 1983, Reports of Progress in Physics, 46, 973
Eenens, P. R. J. & Williams, P. M. 1994, MNRAS, 269, 1082
Eichler, D. & Usov, V. 1993, ApJ, 402, 271
Jones, F. C. & Ellison, D. C. 1991, Space Science Reviews, 58, 259
Marchenko, S. V., Moffat, A. F. J., Ballereau, D. et al. 2003, ApJ, 596, 1295
Monnier, J. D., Traub, W. A., Schloerb, F. P. et al.2004, ApJ, 602, L57
Monnier, J. D., Tuthill, P. G., & Danchi, W. C. 2002, ApJ, 567, L137
Repolust, T., Puls, J., & Herrero, A. 2004, A&A, 415, 349
Setia Gunawan, D. Y. A., van der Hucht, K. A., Williams, P. M., et al. 2001, A&A, 376, 460
Vacca, W. D., Garmany, C. D., & Shull, J. M. 1996, ApJ, 460, 914
White, R. L. & Becker, R. H. 1995, ApJ, 451, 352
Williams, P. M., van der Hucht, K. A., Pollock, A. M. T., et al. 1990, MNRAS, 243, 662
Zhekov, S. A. & Skinner, S. L. 2000, ApJ, 538, 808