1. X-ray Emission from Massive Stars David Cohen Swarthmore College

2. Outline 1.Young OB stars produce strong hard X- rays in their magnetically channeled winds 2.After ~1 Myr X-ray emission is weaker and softer: embedded wind shocks in early O supergiants 3.Line profiles provide evidence of low mass-loss rates

3. Dynamic and energetic processes in outer atmospheres Trace evolution of magnetic fields Provide diagnostics of wind conditions (test theories, fundamental parameters) Influence wind ionization conditions Brightest sources of stellar x-rays - Irradiate circumstellar environment, including nearby systems Relation to diffuse x-rays from bubbles

6.  1 Ori C The red arrow denotes Fe XXV (formed in ~50 MK plasma), which is seen in  1 Ori C but not in  Pup. The blue arrows (light and dark) indicate Si XIV (H-like) and Si XIII (He-like), respectively; note the very different ratios in the two stars. Pup

15. Overall X-ray flux synthesized from the same MHD simulation snapshot. The dip at oblique viewing angles is due to stellar occultation. Data from four different Chandra observations is superimposed. The amount of occultation seen at large viewing angles constrains the radii at which the x-ray emitting plasma exists

16. He-like f/i line ratios in O stars are diagnostics of source location f fi i

18.  1 Ori C The red arrow denotes Fe XXV (formed in ~50 MK plasma), which is seen in  1 Ori C but not in  Pup. The blue arrows (light and dark) indicate Si XIV (H-like) and Si XIII (He-like), respectively; note the very different ratios in the two stars. Pup

19. Differential emission measure: Overall level and temperature distribution of hot plasma are well reproduced by the MHD simulations. On the right is a figure from Gagne et al. (2005) for  1 Ori C. Note the good agreement with the overall shape inferred by W&S from the data. DEM calculated from snapshot of 2-D MHD simulation

21. Some hot stars have x-ray spectra with quite narrow lines, that are especially strong and high energy - not consistent with line-force instability wind shocks  Pup  1 Ori C (O7 V) Capella  1 Ori C is the young hot star at the center of the Orion nebula

23. The line-driven instability (LDI) should lead to shock-heating and X-ray emission 1-D rad-hydro simulation of the LDI

25. But the emission lines are quite broad Ne X Ne IX Fe XVII  Pup (O4 I) 12 Å 15 Å Capella (G 5 III) - a coronal source of soft X-rays

26. Each individual line (here is Ne X Lyat 12.13 Å) is significantly Doppler broadened and blue shifted  Pup (O4 I) Capella (G5 III) HWHM ~ 1000 km/s lab/rest wavelength unresolved at MEG resolution

28. To analyze data, we need a simple, empirical model Detailed numerical model with lots of structure Smooth wind; two- component emission and absorption

29. continuum absorption in the bulk wind preferentially absorbs red shifted photons from the far side of the wind Contours of constant optical depth (observer is on the left) wavelength red blue

30. The basic smooth wind model: for r>R o R o =1.5 R o =3 R o =10 =1,2,8=1,2,8 key parameters: R o &  * j   2 for r/R * > R o, = 0 otherwise

31. Highest S/N line in the  Pup Chandra spectrum o -v ∞ +v ∞ Fe XVII @ 15.014 Å  Fe +16 – neon-like; dominant stage of iron at T ~ 3 X 10 6 K in this coronal plasma 560 total counts note Poisson error bars

32.  * =2.0 R o =1.5 C = 98.5 for 103 degrees of freedom: P = 19%

33. 95% 90% 68% 1.5 <  * < 2.6 and 1.3 < R o < 1.7 1/R o

34. Onset of shock-induced structure: R o ~ 1.5

35. A factor of 4 reduction in mass-loss rate over the literature value of 2.4 X 10 -6 M sun /yr  ~ 150 cm 2 g -1 @ 15 Å 7 X 10 -7 M sun /yr for  * =2

36. Best-fit smooth-wind model with  * = 8 This is the value of  * expected from M = 2.4 X 10 -6 M sun /yr

37. C = 98.5 C = 178 The best-fit model, with  * = 2, is preferred over the  * = 8 model with >99.999% confidence

38. The key parameter is the porosity length, h = (L 3 / l 2 ) = l /f The porosity associated with a distribution of optically thick clumps acts to reduce the effective opacity of the wind h=h’r/R * l ’=0.1 Porosity reduces the effective wind optical depth once h becomes comparable to r/R *

39. h = (L 3 / l 2 ) = l /f

40. The optical depth integral is modified according to the clumping-induced effective opacity: from Owocki & Cohen 2006, ApJ, 648, 565

41. Fitting models that include porosity from spherical clumps in a beta-law distribution: h=h ∞ (1-R * /r)   * =2.0 R o =1.5 h ∞ =0.0 Identical to the smooth wind fit: h ∞ = 0 is the preferred value of h ∞.

42. 95% 68% Joint constraints on  * and h ∞ best-fit model best-fit model with  * =8 C=9.4: best-fit model is preferred over  * =8 model with > 99% confidence

43.  * =8; h ∞ =3.3  * =2; h ∞ =0.0 The differences between the models are subtle… …but statistically significant

44. Two models from previous slide, but with perfect resolution  * =8; h ∞ =3.3  * =2; h ∞ =0.0

45. 95% 68% Joint constraints on  * and h ∞ Even a model with h ∞ =1 only allows for a slightly larger  * and, hence, mass-loss rate h ∞ > 2.5 is required if you want to “rescue” the literature mass-loss rate

46. This degree of porosity is not expected from the line-driven instability. The clumping in 2-D simulations (below) is on quite small scales. Dessart & Owocki 2003, A&A, 406, L1

49. EXTRA SLIDES