CWB logo Configurations in 2008-09 for observers
Colliding-wind binaries
WR140 radio variation
X-MEGA Campaign
Episodic dust-makers
WR140 orbital motion
VdS WR140 campaign
WR140 Introduction
Co-ordinates J2000
R.A. 20 20 27.98
Dec. +43 51 16.3

V = 6.87; v = 7.07

WR 140 (= HD 193793 = BD +43° 3571) is the prototypical Wolf-Rayet plus O star colliding-wind binary which shows periodic, phase-locked variations at all wavelengths, [1] particularly in its X-ray, radio and infrared emission. WR 140 is a remarkable laboratory for a variety of physical processes and, because it is a binary system, the processes repeat each cycle so observers can plan their campaigns to catch particular phenomena. The orbit has high eccentricity and the most extreme wind-collision activity occurs around the time of periastron, when the the separation of the WR and O stars is about 16 times smaller than that at apastron. This occurs for a small fraction of the period. Indeed, the orbital eccentricity is so high that the system gets round three-quarters of its orbit in only 9 months - less than one-tenth of the 8-year period. This is illustrated in the table, which gives phases and dates of orbital phenomena in the months immediately before and after the 2009 periastron passage as well as true anomaly (f), separation of the stars (r/a), and radial velocity of the O star (RV/K), using the orbital elements of Marchenko et al.. Also, Dougherty et al. have determined the inclination (i) and longitude of ascending node, so the orbit can be oriented on the sky. The table gives the position angle of the O star on the sky relative to the WC star and its approximate cardinal point for each phase and also the angles between the WC-O radius vector and the line-of-sight ψ (psi), and the line-of-sight projected on the orbital plane ψp (psip). The angles are shown in the sketch below. The final column gives the expected radial velocity of the compressed Wolf-Rayet wind flowing out along the surface of the wind-collision region (see the figure). The shape of this region is determined by the dynamic balance of the stellar winds and wraps around the O star because its wind is much less dense than that of the WR star. The shape, approximated by a cone whose opening angle depends on the ratio of the momenta of the two stellar winds, is invariant with stellar separation unless the stars get so close that one or both stellar winds suffers radiative braking. The velocity of the compressed flow, vf, depends on the winds of the two stars and the geometry of the interaction region. The radial velocities, RVf1 and RVvf2, of compressed wind observed as emission sub-peaks on top of some He I and C III line profiles vary as the wind-collision structure swings round with the orbital motion, and are given below for two possible values of the cone opening angle: 31° and 50°.

Orbital and wind phenomena at critical phases in 2008-09

phase Year Date Orbital phenomenon f r/a RV/K P.A. pos ψp ψ RVf1 RVf2
0.947 2008.61 Aug 12 conjunction: WC star behind 223 0.62 0.00 83 E 0 32 -1780 -1275
0.996 2009.00 Jan 0 quadrature 314 0.14 +1.00 353 N 90 90 -40 -30
0.000 2009.03 Jan 12 periastron passage 0 0.12 +0.69 323 NW 137 128 +1190 +930
0.004 2009.07 Jan 24 conjunction: O star behind 44 0.14 0.00 262 W 180 148 +1780 +1275
0.046 2009.40 May 26 quadrature 133 0.57 -1.00 173 S 270 90 +30 +22

Observing campaign

The next periastron passage is in January 2009 and the system will be very active from about August 2008 to May 2009. There are plans to observe WR 140 intensively with many instruments in this period: some X-ray observing time has already been awarded and other programmes have started. But WR 140 will be varying so quickly in this period that more observers will be welcome. Fortunately, WR 140 is sufficiently bright to be accessible to many moderate-sized telescopes. Here are some ideas:  

  • Optical spectroscopy to help refine the RV orbit. As can be seen from the table, the velocity changes very rapidly in the second half of 2008 (K(O) = 30 km/s, K(WR) = 82 km/s) and in the first few months of 2009, so frequent - daily - observations are needed to tie down the orbit tightly enough to model the system and all the other phenomena.

  • High-resolution spectroscopy of emission-line profiles to measure the velocity of the compressed wind in the interaction area (blue arrows in the figure), most easily observed as sub-peaks on top of the flat-topped profiles of low-excitation lines like C III at 5696Å [2] or He I at 10830Å [4]. As can be seen from the last column of the table, the range of velocity is enormous. The exact values depend on the geometry of the interaction region.

  • Infrared observation of dust emission. The near-IR flux is expected to rise from quiescence (e.g. K = 5.12) in early January to maximum (K = 3.45) around March 11. The speed of this rise means that it has not been well defined, and frequent observations are needed to define it and the shape of the near-infrared maximum. At longer longer wavelengths, the amplitudes are greater, and the maxima occur later.

  • Optical photometry to observe line-of-sight dust. Some of the dust formed by WR 140 will lie in our line of sight, causing small eclipses in the optical. The colours of the eclipses give information on the dust grain sizes [2], and frequent, multi-colour optical monitoring from December to May would be valuable. 

See also the VdS Spectroscopy Sections's WR 140 campaign pages.


Below: sketch of the WR 140 system illustrating the angles in the table. Only the orbit of the O star is shown (green, dotted line). As the inclination is about 122°, we view the system from below the orbital plane, and the O star moves clockwise on the sky. The thick grey line represents the wind-collision zone, approximately a cone at large distances from the stars, and compressed stellar wind flows along it with velocity vf determined by the winds of the two stars and the geometry of the wind-collision zone.
rev: 13 March 2008
Peredur Williams
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[1] P.M. Williams et al. MNRAS 243, 662, 1990
[2] S.V Marchenko et al. ApJ 596, 1295, 2003.
[3] S.M. Dougherty et al. ApJ 623, 447, 2005.
[4] W.P. Varricatt et al. MNRAS 351, 1307, 2004