Chapter 1 Introduction

Chapter 1 Introduction 1.1 Why study massive stars? Massive, early-type stars (M > 10M ), despite their extreme rareness (for every 20M star in ∼ the Milky Way there are 105 solar-type stars; Massey 2003), have profound effects on their ∼ environments. They are the main source of ultraviolet (UV) ionizing radiation in galaxies, and through the heating of dust they power galaxies' far-infrared (FIR) luminosities. Their powerful stellar winds (see below) and their violent deaths as supernovae provide most of mechanical en- ergy input into the interstellar medium (ISM; Abbott 1982a). They are important for galactic chemical evolution, as they are the main source of carbon, nitrogen and oxygen (CNO) enrich- ment of the ISM (Maeder 1981). They are also the only places in the universe where elements more massive than iron can be synthesized, either through the slow capture of neutrons (the s process) chiefly during helium burning, or through the rapid capture of neutrons (the r process) in the core-collapse supernova explosion (Woosley et al. 2002). The collective action of a group of massive stars formed together in a rapid \starburst" can give rise to a galactic-scale super- wind (Heckman 1998) which chemically enriches (and may also heat) the intergalactic medium (Donahue 1998). As well as these environmental effects, the core collapse of more massive stars (M > 25M ) probably leads to the formation of stellar-mass black holes (Horowitz & Teukolsky ∼ 1999). Finally, massive star supernovae are widely believed to be connected with gamma-ray bursts { among the most energetic phenomena in the Universe (Kaper 2003; Fryer 2003). Massive stars are therefore of great importance to many areas of astrophysics. In this thesis we concentrate on one particular aspect of massive stars: their powerful winds, and in particular, high-resolution spectroscopy of the X-rays produced when the winds of two such stars in a binary system collide. In the remainder of this introduction, we first give an overview of massive star winds in Section 1.2. In Section 1.3 we describe X-rays from single 2 Chapter 1. Introduction 3 massive stars. In Section 1.4 we describe the X-rays produced in a wind-wind collision in a so-called colliding wind binary (CWB), and we describe observations of these systems at other wavelengths in Section 1.5. After a brief diversion to discuss wind-wind collisions in other systems in Section 1.6, we give a brief overview of X-ray telescopes and X-ray spectroscopy in Section 1.7, and finally introduce the remainder of the thesis in Section 1.8. 1.2 Mass-loss from massive stars 1.2.1 The discovery of massive stars winds As mentioned above, hot, massive, early-type stars { Wolf-Rayet, O-type and early B-type (B0 1 −7 −5 −1 or B1) stars { have powerful winds, with mass-loss rates M_ 10 {10 M yr and terminal ∼ −1 velocities v1 1000{3000 km s (e.g. Corcoran et al. 2004; Pittard et al. 2004). Cooler, later ∼ stars may have comparable mass-loss rates (e.g. some red supergiants), but stars later than B1 (cooler than 21 000 K) have much lower wind speeds (e.g. a few hundred km s−1 for late B stars, < 100 km s−1 for red supergiants and 10 km s−1 for asymptotic giant branch stars; Lamers ∼ ∼ −14 −1 & Cassinelli 1999). For comparison, the mass-loss rate of the Sun is 10 M yr (e.g. ∼ Cassinelli 2002) and its wind speed is a few hundred km s−1 (e.g. Holzer 1989). The mechanical 2 power (M_ v1=2) of a massive star's wind is thus roughly 9 or 10 orders of magnitude larger than that of the Sun. The discovery of mass-loss from stars dates back to the nineteenth century, when it was found that the optical spectrum of P Cygni showed broad emission lines with absorption on their blueward edges (e.g. Bélopolsky 1899). Such profiles were seen in a number of other stars (e.g. Campbell 1895; Beals 1929, 1931). Beals (1929) in particular noted the phenomenon in Wolf-Rayet stars (described in more detail in Section 1.2.2 below). He interpreted these profiles as evidence of outflowing mass from the star. The emission component comes from the scattering of photospheric light into the line of sight from all parts of the stellar wind around the star (except for the small portion that is occulted by the stellar disk). This gives rise to emission with Doppler velocities between v1 and +v1 { the exact shape of the emission − ≈ component depends on the velocity law v(r) (i.e. on how the wind speed v varies with distance r 1 Due to their similarities, such as high Teff and powerful winds, O- and early B-type stars are often collectively referred to as OB stars. This is particularly used to describe an early-type star whose precise spectral classification is uncertain. Chapter 1. Introduction 4 from the centre of the star). The absorption component arises in the portion of the wind in front of the stellar disk, as photospheric light is scattered out of the line of sight { this component has Doppler velocities ranging from v1 to 0. The sum of these two contributions gives rise to a P − Cygni profile (see Fig. 1.1). From the blueward edge of the absorption trough one may measure v1; from his stellar spectra Beals (1929) found outflow speeds ranging from a few hundred to 2400 km s−1. Later, Wilson (1958) noted broad emission lines in the optical spectra of O stars, similar to those described by Beals (1929). He interpreted these as evidence that O stars too are losing mass, with wind speeds of 1500 km s−1. ∼ The study of mass-loss from early-type stars took a great leap forward in the 1960s with the beginning of space-based observations. This was pioneered by Donald Morton and his co- workers, who used rocket-borne UV spectrographs to measure the spectra of Orion OB stars (Morton 1967a,b; Morton et al. 1968). From the P Cygni profiles in the spectra of three of these −6 −1 stars (δ, and ζ Orionis), Morton (1967b) inferred mass-loss rates of 10 M yr and wind ∼ speeds of 1500 km s−1. These early studies led to a great increase in interest in massive star ∼ winds from the late sixties onwards. One of the first questions to be addressed was how the winds are accelerated up to such high speeds with such high mass-loss rates. 1.2.2 The radiative driving of massive star winds Lucy & Solomon (1970) showed that radiation pressure on the very same UV P Cygni lines which were used to infer mass loss could accelerate the winds up to the high observed speeds −1 −8 −1 (up to 3000 km s ). However, their model predicted mass-loss rates of 10 M yr , 1 or ∼ ∼ 2 orders of magnitude lower than those that are observed. This is because they only included the effects of relatively few optically thick resonance lines. Castor, Abbott, & Klein (CAK; 1975) developed a powerful method for including the effects of many lines, optically thick and thin, resonant and non-resonant. In particular, they introduced the so-called CAK force multiplier M(t) for expressing the total radiative acceleration gL due to all spectral lines: σ L g = e ? M(t); (1.1) L 4πr2c 2 −1 where σe is the electron scattering opacity (0:34 cm g for solar-abundance material; Bowers & Deeming 1984), L? is the stellar luminosity, r is the distance from the centre of the star, and Chapter 1. Introduction 5 Figure 1.1: How a P Cygni profile originates in a stellar wind. The `halo' (H) surrounding the star gives rise to the emission component. The `tube' in front of the star (F) gives rise to the absorption component, while the `tube' behind the star (O) is occulted. The bottom panel shows the resulting P Cygni profile. (Figure from Lamers & Cassinelli 1999.) Chapter 1. Introduction 6 c is the speed of light. CAK showed that the force multiplier could be expressed as a power law: M(t) = kt−α; (1.2) where k is a normalizing constant, and t is a dimensionless optical depth parameter, given by dv −1 t = σ v ρ ; (1.3) e T dr where vT is the thermal velocity, ρ is the density and v is the wind speed. The exponent α in equation (1.2) depends on the relative contributions of optically thick and thin lines: if all lines are thick α = 1, whereas if they are all thin α = 0. CAK found α 0:7{0.8. Note ≈ that the line driving force depends upon the velocity gradient dv=dr. When there is a velocity gradient, higher, faster-moving layers of the stellar wind have their lines Doppler-shifted out of the `shadow' of the lower layers. This means that the line driving force can persist out to large radii. CAK found the line driving force to be 100 times larger than Lucy & Solomon's (1970) ∼ value, and therefore showed that line driving could explain the high mass-loss rates observed from early-type stars, as well as the high wind speeds. CAK theory predicts that the terminal velocity should scale with the escape velocity at the stellar surface vesc: α v1 = v ; (1.4) 1 α esc r − and that the velocity should vary with r as follows: 0:5 R∗ v(r) = v1 1 ; (1.5) r − where R∗ is the stellar radius.

Load more