The diversity of High-Mass X-ray binaries

1. The diversity of High-Mass X-ray binaries Agios Nikolaos October 2010 Ignacio Negueruela where astrologers roam …

2. High-mass X-ray binaries • High-mass X-ray binaries are systems containing a compact object accreting from a massive star. • Good separation from LMXBs. • Very few intermediate-mass objects: LMC X-3, V4641 Sgr • Fundamental tools for astrophysics, cf. Cen X-3 & Cyg X-1. • Massive star is the main (only) contributor to optical and infrared brightness.

3. X-ray pulsars • Most massive X-ray binaries are X-ray pulsars magnetised neutron stars • We know of two black hole systems, one in the Milky Way (Cyg X-1; O9.7Iab), and one in the LMC (LMC X-1, O8III-V) . • Others are being found in nearby galaxies: M33 X-7 and IC 10 X-1 (talks by Fabbiano, Pietsch) • Weird cases: Cyg X-3, SS433; HMXBs? Related to ULXs? (talk by Roberts) • Presence of strong magnetic field somehow inhibits jet formation no radio detections. • Nature of donor determines main properties.

4. X-ray spectra of accreting pulsars • The X-ray spectra of accreting pulsars are generally fitted to phenomenological models (power-law+cutoff) • Increasing effort to interpret them in physical terms. Bulk Comptonisation of thermal components (e.g., Becker & Wolff 2007, ApJ 654, 435) Talk by Haberl

5. X-ray pulsars Be star Supergiant Modern version of Corbet’s diagram (Corbet 1986, MNRAS 220, 1047)

6. Classes of HMXBs Be/X-ray binaries Accretion from the wind of a supergiant Roche-lobe overflow

7. Classical HMXB Cen X-3 SMC X-1 + LMC X-1 (BH) LMC X-4

8. Classical HMXBs • Short orbital periods (2-3 days) • Circularised orbits • Incipient Roche-lobe overflow • The stars may be bloated, and are over-luminous for their mass • Formation of an accretion disk results in high LX  detectable in other galaxies Van der Meer et al. (2007, A&A 473, 523)

9. Classical HMXBIncipient Roche-lobe overflow LX 1038 erg s-1

10. Formation channel for HMXBs • Case C mass transfer • q << 1 • Non-conservative evolution via common envelope • Results in SG+BH • Only way to make a BH in a binary? (talk by Casares) Wellstein & Langer (1999; A&A 350, 148)

11. Note: many more Be/X in the SMC Talk by Coe Be/X-ray binaries

12. Be/X-ray binaries a Be star isan early type (O7 to A1) star, not very evolved (luminosity class III-V), which shows - or has shown - emission in the H line (see Porter & Rivinius 2003, PASP 115, 1153 for a review). • Other Balmer and singly-ionised metallic lines (Fe II, Cr II, etc) also seen in emission. HeI in emission in stars earlier than B2. At sufficient resolution, all lines are double peaked. • There is also an infrared excess due to continuum emission.

13. Be/X-ray binaries • These characteristics can be explained by the presence of a disk of material expelled from the star. • Currently the model favoured is the decretion viscous disk (Lee et al. 1991, MNRAS 250, 432), which can reproduce most observational characteristics (Porter 1999, A&A 348, 512; Okazaki 2001, PASJ 53, 119) • At a given time, around 10% of early-B stars are in a Be star phase • But the Be phenomenon is very variable. Stars move from Be to non–Be phase.

14. Therefore … A Be/X-ray binary is made of • A Be star – observationally always O9-B1 both in the Galaxy and the LMC • A compact object accreting material from the disk of the Be star – observationally always a neutron star ►X-ray pulses detected whenever one looks hard enough. Indistinguishable distribution in the SMC (McBride et al. 2008; MNRAS 388, 1198)

15. X-ray lightcurves Persistent sources • Relatively low LX ( 1034 erg s-1). • Small intensity fluctuations (factor  10) without an obvious temporal pattern. Transients • Quiescence: low (≤1035 erg s-1) or non-detectable LX. • Series of outbursts with relatively high X-ray luminosity (Lx  1037 erg s-1), separated by the (suspected) orbital period (Type I or normal according to Stella et al. 1986, ApJ 308, 669). • Larger outbursts with Lx> 1037 erg s-1 (Lx ≈ LEdd ), lasting several weeks and not showing modulation with the orbital period (giant or Type II) .

16. X-ray lightcurve of the persistent Be/X-ray binary X Per (4U 0352+30), taken with the All Sky Monitor on board RossiXTE Porb = 250.0 d, Pspin = 837.6 s, e = 0.11

17. X-ray lightcurves Persistent sources • Relatively low LX ( 1034 erg s-1). • Small intensity fluctuations (factor  10) without an obvious temporal pattern. Transients • Quiescence: low (≤1035 erg s-1) or non-detectable LX. • Series of outbursts with relatively high X-ray luminosity (Lx  1037 erg s-1), separated by the (suspected) orbital period (Type I or normal according to Stella et al. 1986, ApJ 308, 669). • Larger outbursts with Lx> 1037 erg s-1 (Lx ≈ LEdd ), lasting several weeks and not showing modulation with the orbital period (giant or Type II) .

18. X-ray lightcurve of the prototype Be/X-ray transient EXO 2030+375, taken with the All Sky Monitor on board RossiXTE Porb = 46.0 d, Pspin = 41.7s, e = 0.41

19. X-ray lightcurve of the Be/X-ray transient MXB 0656 -072, taken with the All Sky Monitor on board RossiXTE. The only previous recorded outburst took place in 1974 (but there was another one four years later). Porb = ? d, Pspin = 160.7s

20. Several transients display series of Type I outbursts after (and only after) a giant outburst. 2S 1417-624Porb = 42.1 d, Pspin= 17.6s, e = 0.45

21. Optical studies • Optical monitoring reveals strong changes in the line profiles  tracers of the disk’s dynamics • These changes are generally accompanied by large photometric variability • They can be explained as large variations in the disk’s configuration Reig et al. (2007, A&A 462, 1081)

22. The truncated disk modelOkazaki & Negueruela (2001, A&A 377, 161) • The phenomenology observed implies strong interaction between the different system components. This can be easily understood in terms of the decretion disk model. • If the disk is supported by viscosity, the neutron star exerts a torque on disk particles that makes them lose angular momentum. • As a consequence, the disk can only grow up to a certain size, and will be truncated at one of the commensurabilities between the Keplerian orbital period of the neutron star and disk particles the disk acts as reservoir of mass.

23. Model successes The model effectively explains two observational facts: • There is a good correlation between the orbital period and the maximum EW(H) measured (Reig et al. 1997, A&A 322, 193)  the neutron star controls the size of the disk. • Analysis of emission-line shapes and infrared excess indicates rather higher densities in the disks of Be/X-ray binaries than in those of isolated Be stars (Zamanov et al. 2001, A&A 367, 884). The model predicts a strong dependence of the observed behaviour on the orbital eccentricity, which is generally observed to hold.

24. But there are exceptions … • The more Be/X-ray binaries we know, the more difficult it seems to find a common pattern in their behaviour. • I trust, however, that all those different behaviours arise from the very complex dynamical interplay between the components of the systems and can finally be reduced to the same physical processes. KS 1947+300 Porb= 40.4 d, Pspin= 18.7s, e = 0.03

25. Where do they come from? • Be/X-ray binaries are descended from moderately massive binaries that undergo a phase of mass transfer • They are believed to originate from relatively close systems with q<0.5 in which semi-conservative mass transfer is possible • Typical age ≥ 10 Myr

26. See Coe’s talk for SMC population Considered as a population, BeXBs can be used to set constraints on formation models and hence on basic physics. • There is growing evidence that a substantial population of Be/X-ray binaries with low eccentricity exists. • Inference of electron-capture SN  dependence on previous binary history. Podsiadlowski et al. 2004, ApJ 612, 1044

27. The Be + WD mystery • Population synthesis models provide tools to analyse populations (e.g., Van Bever & Vanbeveren 1997, A&A 322, 116; Raguzova 2001 A&A 367, 848). • All population synthesis models that have been elaborated predict that, for every Be + neutron star binary, there should be ~ 10 Be + WD binaries. • No such system has been conclusively identified. They are very hard to pinpoint, but there should be many!

28. The Be + WD mystery • Population synthesis models provide tools to analyse populations (e.g., Van Bever & Vanbeveren 1997, A&A 322, 116; Raguzova 2001 A&A 367, 848). • All population synthesis models that have been elaborated predict that, for every Be + neutron star binary, there should be ~ 10 Be + WD binaries. • No such system has been conclusively identified. They are very hard to pinpoint, but there should be many! This renders the models somewhat suspect!

29. There’s a correlation here … Be/X-ray binaries

30. Equilibrium at which corotation velocity at the magnetospheric radius equals Keplerian velocity (Corbet 1986, MNRAS 220, 1047; Waters & van Kerkwijk 1989, A&A 223, 196)

31. Equilibrium at which corotation velocity at the magnetospheric radius equals Keplerian velocity (Corbet 1986, MNRAS 220, 1047; Waters & van Kerkwijk 1989, A&A 223, 196)

32. Evolved O8-B2 stars (luminosity class I)

33. Supergiant X-ray binaries Vela X-1: • Short term flaring • Long term variability by a factor of 4 Ribó et al. 2006 (A&A, 449, 687) Flare from 4U 1907+09Fritz et al. 2006 (A&A 458, 885)

34. Supergiant X-ray binaries

35. Radiative winds from hot stars Line Scattering: Bound Electron Resonance Heavy ions have large Thompson cross sections The  law   0.8 – 1.2 Review: Kudritzki & Puls 2000, ARA&A, 38, 613 Images stolen from Stan Owocki

36. Velocity Density Development of instability smooth wind Owocki & Rybicki 1984, ApJ, 284, 337 cf. Feldmeier et al. 1997, A&A, 322, 878 Images stolen from Stan Owocki

37. Wind clumping • Clumping factor • Size and geometry of clumps • Shells or blobs • Optically thin? 1D simulations Runacres & Owocki 2002, A&A, 381, 1015 2D simulations Dessart & Owocki 2003, A&A, 406, L1 Porous winds Owocki et al. 2004, ApJ, 616, 525 Oskinova et al. 2006, MNRAS, 372, 313 Constraints from spectra Prinja et al. 2005, A&A 430, L41 Bouret et al. 2005, A&A, 438, 301 Puls et al. 2006, A&A, 454, 625

38. Bondi-Hoyle-Lyttleton accretion Dependence of LX on eccentricity Reig et al. (2003, A&A 405, 285) See review: Edgar 2004, New Ast. Rev. 48, 843

39. But this also becomes unstable … Transverse instability close to stagnation point (Foglizzo et al. 2005; A&A 435, 397) Can (transient) accretion disks form? Perturbed accretion flow (Frixell & Taam 1988, ApJ 335, 862)

40. This is a complex problem A photo-ionization wake in Vela X-1 Kaper et al. (1994, A&A 289, 846)

41. This is a complex problem A photo-ionization wake in Vela X-1 Kaper et al. (1994, A&A 289, 846) A tidally induced accretion stream forms in the models of Blondin et al. (1991, ApJ 371, 684), which incorporate a realistic representation of the physics.

42. X-rays and winds A photo-ionization wake in Vela X-1 Kaper et al. (1994, A&A 289, 846)

43. There is feedback everywhere Tidally induced non-radial pulsations in Vela X-1 Quaintrell et al. (2003, A&A 401, 313)

44. Formation channel for SGXBs • Case A mass transfer • q 1 • Conservative evolution with two mass-transfer phases • Results in SG+NS Wellstein & Langer (1999; A&A 350, 148)

45. Supergiant Fast X-ray transients • A group of flaring sources with very short outbursts and supergiant companions. • Transient emission composed by many flares reaching LX 1036 -1037 erg s-1 • Persistent emission at lower luminosity LX 1033 -1034 erg s-1 • Deep quiescence at LX 1032 erg s-1 (Giunta et al. 2009, MNRAS 399, 744; Bozzo et al. arXiv:1004.2059; Sidoli et al. 2010 arXiv:1007.1091) e.g., Romano et al. 2010, Mem. SAI 81, 332

46. Parameters of SFXTs Optical counterpart to AX 1845.0-0433 (VLT+FORS1)

47. The real outbursts Three days of Suzaku observations of IGR J17544-2619 • Outburst (series of flares) is  1 day. • Single flare is  3 minutes • Fastest doubling time is  4 seconds. Data and graphics by courtesy of D. M. Smith

48. Looking for a difference • This leaves several options: • Difference in wind structure • Difference in wind geometry • Difference in accretion process

49. Clumpy wind • First proposed by in’t Zand (2005, A&A 441, L1) to explain behaviour of IGR J17544-2619. • All winds from OB stars are likely clumpy. • Classical supergiant X-ray binaries also show flares. INTEGRAL monitoring of Vela X-1 (Kreykenbohm et al. 2008, A&A 492, 511)

50. Equatorial overdensity • Model proposed by Sidoli et al. (2007, A&A 476, 1307) to explain behaviour of IGR J11215-5952. • Not clear how to extend it to other sources. • Are winds spherically symmetric?