X-ray Emission from Massive Stars

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

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

3. Role of X-rays in OB Stars Trace 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 OB stars are brightest sources of stellar x-rays - irradiate circumstellar environment, including nearby stars, disks, planetary atmospheres Relation to diffuse x-rays from bubbles?

4. Young OB stars are very X-ray bright Lx: up to ~1034 ergs s-1 X-ray temperatures: few up to 10+ keV (10s to 100+ million K) Orion; COUP (Feigelson et al. 2002)

5. ~17’

7. 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) indicateSi XIV (H-like)andSi XIII (He-like), respectively; note the very different ratios in the two stars. 1 Ori C z Pup

16. 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

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

19. 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) indicateSi XIV (H-like)andSi XIII (He-like), respectively; note the very different ratios in the two stars. 1 Ori C z Pup

20. 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 q1 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

22. 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 q1Ori C(O7 V) Capella zPup q1Ori C is the young hot star at the center of the Orion nebula

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

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

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

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

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

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

32. Highest S/N line in the z Pup Chandra spectrum Fe XVII @ 15.014 Å -v∞ lo +v∞ 560 total counts note Poisson error bars l/Dl Fe+16 – neon-like; dominant stage of iron at T ~ 3 X 106 K in this coronal plasma

33. *=2.0Ro=1.5 C = 98.5 for 103 degrees of freedom: P = 19%

34. 95% 90% 1/Ro 68% 1.5 < * < 2.6 and 1.3 < Ro < 1.7

35. Onset of shock-induced structure: Ro ~ 1.5

36. for t*=2 ~ 150 cm2 g-1 @ 15 Å 7 X 10-7 Msun/yr A factor of 4 reduction in mass-loss rate over the literature value of 2.4 X 10-6 Msun/yr

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

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

39. The porosity associated with a distribution of optically thick clumpsacts to reduce the effective opacity of the wind h=h’r/R* l’=0.1 The key parameter is the porosity length, h = (L3/l2) = l/f Porosity reduces the effective wind optical depth once h becomes comparable to r/R*

40. h = (L3/l2) = l/f

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

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

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

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

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

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

47. 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

50. EXTRA SLIDES