X-ray Emission from Massive Stars David Cohen Swarthmore College

X-ray Emission from Massive Stars

David Cohen

Swarthmore College

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

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

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

QuickTime™ and aBMP decompressor

are needed to see this picture.

QuickTime™ and aBMP decompressor

are needed to see this picture.

QuickTime™ and aBMP decompressor

are needed to see this picture.

QuickTime™ and aYUV420 codec decompressor

are needed to see this picture.

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

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

ff i

i

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

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

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

Pup1 Ori C

(O7 V)Capella

1 Ori C is the young hot star at the center of the Orion nebula

The line-driven instability (LDI) should lead to shock-heating and X-ray emission

1-D rad-hydro simulation of the LDI

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

Each individual line (here is Ne X Lyat 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

To analyze data, we need a simple, empirical model

Detailed numerical model with lots of structure

Smooth wind; two-component emission

and absorption

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

redblue

The basic smooth wind model:

for r>Ro

Ro=1.5

Ro=3

Ro=10

=1,2,8

€

∗ ≡κM

⋅

4πR∗v∞

μπ λ drdrjeL

R

21

1

2

*

8 ∫ ∫−

∞ −=

key parameters: Ro & *

j 2 for r/R* > Ro,

= 0 otherwise

€

=∗R∗dz'

r'2 (1− R∗r ')β

z

∞

∫

Highest S/N line in the Pup Chandra spectrum

λo-v∞ +v∞

Fe XVII @ 15.014 Å

λλ

Fe+16 – neon-like; dominant stage of iron at T ~ 3 X 106 K in this coronal plasma

560 total counts

note Poisson error bars

*=2.0Ro=1.5

C = 98.5 for 103 degrees of freedom: P = 19%

95%

90%

68%

1.5 < * < 2.6 and 1.3 < Ro < 1.7

1/Ro

Onset of shock-induced structure: Ro ~ 1.5

A factor of 4 reduction in mass-loss rate over the literature value of 2.4 X 10-6 Msun/yr

€

∗ =3.6κ150 M−6

⋅

R12v2000

M•

−6 =τ ∗R12v2000

3.6κ150

€

∗ ≡κM

⋅

4πR∗v∞

κ ~ 150 cm2 g-1 @ 15 Å

7 X 10-7 Msun/yr

for *=2

Best-fit smooth-wind model with * = 8

This is the value of * expected from M = 2.4 X 10-6 Msun/yr

C = 98.5C = 178

The best-fit model, with * = 2, is preferred over the * = 8 model with

>99.999% confidence

The key parameter is the porosity length, h = (L3/l2) = 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*

h = (L3/l2) = l/f

The optical depth integral is modified according to the clumping-induced effective opacity:

€

κeff =κ 1− e−τ c( )

τ c

from Owocki & Cohen 2006, ApJ, 648, 565

Fitting models that include porosity from spherical clumps in a beta-law distribution:

h=h∞(1-R*/r)

*=2.0Ro=1.5h∞=0.0

Identical to the smooth wind fit: h∞ = 0 is the preferred value of h∞.

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

*=8; h∞=3.3*=2; h∞=0.0

The differences between the models are subtle…

…but statistically significant

Two models from previous slide, but with perfect resolution

*=8; h∞=3.3*=2; h∞=0.0

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

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

EXTRA SLIDES