x-ray emission and absorption in massive star winds
DESCRIPTION
X-ray Emission and Absorption in Massive Star Winds. David Cohen Swarthmore College. Constraints on shock heating and wind mass-loss rates. What we know about X-rays from single O stars. David Cohen Swarthmore College. via high-resolution spectroscopy. Outline. Morphology of X-ray spectra - PowerPoint PPT PresentationTRANSCRIPT
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X-ray Emission and Absorption in Massive Star
WindsConstraints on shock
heating and wind mass-loss rates
David CohenSwarthmore College
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What we know about X-rays from single O stars
via high-resolution spectroscopy
David CohenSwarthmore College
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OutlineMorphology of X-ray spectra Embedded Wind Shocks: X-ray
diagnostics of kinematics and spatial distribution
Line profiles and mass-loss ratesBroadband absorption
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Morphology Pup (O4 If)
Capella (G5 III) – coronal source – for comparison
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Pup (O4 If)
Capella (G5 III) – coronal source – for comparison
Mg XIMg XII
Si XIIISi XIV Ne IXNe X Fe XVII
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Pup (O4 If)
Capella (G5 III) – coronal source – for comparison
Mg XIMg XII
Si XIIISi XIV Ne IXNe X H-like vs. He-like
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Pup (O4 If)
Capella (G5 III) – coronal source – for comparison
Mg XIMg XII H-like vs. He-like
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Pup (O4 If)
Capella (G5 III) – coronal source – for comparison
Mg XIMg XII H-like vs. He-like
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Spectral energy distribution trendsThe O star has a harder spectrum, but
apparently cooler plasma
We’ll see later on that soft X-ray absorption by the winds of O stars explains this
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NextFirst, we’ll look at what can be learned
about the kinematics and location of the X-ray emitting plasma from the emission lines
After that, we’ll look at the issues of wind absorption and its effects on line shapes and on the broadband hardness trend
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Embedded wind shocksNumerical simulations of the line-
driven instability (LDI) predict:1. Distribution of shock-heated
plasma2. Above an onset radius of r ~
1.5 R*
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1-D rad-hydro simulationsSelf-excited instability (smooth initial
conditions)
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shock onset at r ~ 1.5 RstarVshock ~ 300 km/s :
T ~ 106 K
pre-shock wind plasma has low density
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shock onset at r ~ 1.5 RstarVshock ~ 300 km/s :
T ~ 106 K
Shocked wind plasma is decelerated back down to the local CAK wind velocity
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Morphology Pup (O4 If)
Capella (G5 III) – coronal source – for comparison
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Morphology – line widths Pup (O4 If)
Capella (G5 III) – coronal source – for comparison
Ne X Ne IX Fe XVII
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Morphology – line widths Pup (O4 If)
Capella (G5 III) – coronal source – for comparison
Ne X Ne IX Fe XVII
~2000 km/s ~ vinf
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Shock heated wind plasma is moving at >1000 km/s :
broad X-ray emission lines
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99% of the wind mass is cold*, partially ionized…
x-ray absorbing
*typically 20,000 – 30,000 K; maybe better described as “warm”
opacity
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x-ray absorption is due to bound-free opacity of
metals
…and it takes place in the 99% of the wind that is unshocked
opacity
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Emission + absorption = profile modelThe kinematics of the emitting
material dictates the line width and overall profile
Absorption affects line shapes
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observer on left
isovelocity contours
-0.8vinf
-0.6vinf
-0.2vinf +0.2vi
nf
+0.6vi
nf
+0.8vi
nf
Onset radius of X-ray emission, Ro
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observer on left
isovelocity contours
-0.8vinf
-0.6vinf
-0.2vinf +0.2vi
nf
+0.6vi
nf
+0.8vi
nf
optical depth
contours
t = 0.3
t = 1
t = 3
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Profile shape assumes beta velocity law and onset radius, Ro
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t tRdz'
r' 2 (1 R r ')z
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t tRdz'
r' 2 (1 R r ')z
Universal property of the wind
z different for each point
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t M
4Rv
opacity of the cold wind wind mass-loss rate
wind terminal velocityradius of the star
€
M•
= 4πr2vρ
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t=1,2,8
key parameters: Ro & t*
t M
4Rv
j ~ 2 for r/R* > Ro
= 0 otherwise
t tRdz'
r' 2 (1 R r ')z
Ro=1.5
Ro=3
Ro=10
t = 1 contours
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We fit these x-ray line profile models to each line in the Chandra data
Fe XVIIFe XVII
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And find a best-fit t* and Ro…
Fe XVIIFe XVII t* = 2.0
Ro = 1.5
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…and place confidence limits on these fitted parameter values
68, 90, 95% confidence limits
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Let’s focus on the Ro parameter first
Note that for = 1, v = 1/3 vinf at 1.5 R* and 1/2 vinf at 2 R*
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Distribution of Ro values in the Chandra spectrum of Pup
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Vinf can be constrained by the line fitting too
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Wind AbsorptionNext, we see how absorption in the
bulk, cool, partially ionized wind component affects the observed X-rays
CNO processed
Solar
opacity
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Morphology Pup (O4 If)
Capella (G5 III) – coronal source – for comparison
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Morphology – line widths Pup (O4 If)
Capella (G5 III) – coronal source – for comparison
Ne X Ne IX Fe XVII
~2000 km/s ~ vinf
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Pup (O4 If)
Capella (G5 III) – unresolved
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t M
4Rv
opacity of the cold wind wind mass-loss rate
wind terminal velocityradius of the star
€
M•
= 4πr2vρ
t* is the key parameter
describing the absorption
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observer on left
optical depth
contours
t* = 2 in this wind
-0.8vinf
-0.6vinf
-0.2vinf +0.2vi
nf
+0.6vi
nf
+0.8vi
nf
t = 0.3
t = 1
t = 3
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Wind opacity X-ray bandpass
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Wind opacity X-ray bandpass
Chandra HETGS
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X-ray opacity CNO processed
Solar
Zoom in
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We fit these x-ray line profile models to each line in the Chandra data
Fe XVIIFe XVII
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And find a best-fit t* and Ro…
Fe XVIIFe XVII t* = 2.0
Ro = 1.5
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…and place confidence limits on these fitted parameter values
68, 90, 95% confidence limits
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Pup: three emission lines
Mg Lya: 8.42 Å Ne Lya: 12.13 Å O Lya: 18.97 Å
t* = 1 t* = 2 t* = 3
Recall:
t M
4Rv
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atomic opacity of the wind CNO processed
Solar
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Results from the 3 line fits shown previously
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Fits to 16 lines in the Chandra spectrum of Pup
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Fits to 16 lines in the Chandra spectrum of Pup
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t M
4Rv
t*(l) trend consistent with (l)
Fits to 16 lines in the Chandra spectrum of Pup
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Fits to 16 lines in the Chandra spectrum of Pup
CNO processed
Solar
t*(l) trend consistent with (l)
t M
4Rv
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t M
4Rv
t*(l) trend consistent with (l)
M becomes the free parameter of the fit to the t*(l) trend
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t M
4Rv
t*(l) trend consistent with (l)
M becomes the free parameter of the fit to the t*(l) trend
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Traditional mass-loss rate: 8.3 X 10-6 Msun/yr
Our best fit: 3.5 X 10-6 Msun/yr
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Fe XVIITraditional mass-loss rate: 8.3 X 10-6 Msun/yr
Our best fit: 3.5 X 10-6 Msun/yr
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Ori: O9.5e Ori: B0
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Ori (09.7 I): O Lya 18.97 Å
Ro = 1.6 R*t* = 0.3
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t* quite low: is resonance scattering affecting this line?
Ro = 1.6 R*t* = 0.3
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e Ori (B0 Ia): Ne Lya 12.13 Å
Ro = 1.5 R*t* = 0.6
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HD 93129Aab (O2 If*): Mg XII Lya 8.42 Å
Ro = 1.8 R*t* = 1.4
M-dot ~ 2 x 10-5 Msun/yr from unclumped Ha
Vinf ~ 3200 km/s
M-dot ~ 5 x 10-6 Msun/yr
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HD93129A – O2 If*Most massive, luminous O star (106.4
Lsun)
Strongest wind of any O star (2 X 10-5 Msun/yr; vinf = 3200 km/s)
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Its X-ray spectrum is hard
low H/He high H/He
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Its X-ray spectrum is hard
low H/He
high H/He
But very little plasma with kT > 8 million K
MgSi
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Low-resolution Chandra CCD spectrum of HD93129AFit: thermal emission with wind + ISM absorption
plus a second thermal component with just ISM
kT = 0.6 keV *wind_abs*ismadd 5% kT = 2.0 keV*ism
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kT = 0.6 keV *wind_abs*ismadd 5% kT = 2.0 keV*ism
Typical of O stars like Pup
t*/ = 0.03 (corresp. ~ 5 x 10-6 Msun/yr)
Small contribution from colliding wind shocks
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What about broadband effects?The X-ray opacity changes across the
Chandra bandbass – this should affect the overall X-ray SED: soft-X-ray absorption in O star winds should harden the emergent X-rays
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Emission measure and optical depth
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Emergent X-ray flux
exact exospheric
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Transmission
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S* is the free parameterOnce you assume a run of opacity vs. wavelength, (l)
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exponential absorption (from a slab) is inaccurate
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Wind opacity X-ray bandpass
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Combining opacity and transmission models
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Broadband trend is mostly due to absorption
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Although line ratios need to be analyzed too
low H/He high H/He
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ConclusionsWind absorption has an important effect on
the X-rays we observe from single O starsResolved line profiles show widths consistent
with shock onset radii of ~1.5 R*
And shapes indicative of wind absorption – though with low mass-loss rates
Initial results on broadband X-ray spectra indicate that wind absorption is significant there, too
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Additional SlidesClumping and por0sity
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Multi-wavelength evidence for lower mass-loss rates 2005 onward
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Pup mass-loss rate < 4.2 x 10-6 Msun/yr
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“Clumping” – or micro-clumping: affects density-squared diagnostics; independent of clump size, just depends on clump density contrast
(or filling factor, f )
visualization: R. Townsend
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The key parameter is the porosity length,
h = (L3/ℓ2) = ℓ/f
But porosity is associated with optically thick clumps, it acts to reduce the effective opacity of the wind; it does depend on the size
scale of the clumps
L
ℓ
f = ℓ3/L3
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h = ℓ/f
clump size
clump filling
factor (<1)
The porosity length, h:
Porosity length increasing
Clump size increasing
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Porosity only affects line profiles if the porosity length (h) exceeds the stellar radius
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The clumping in 2-D simulations (density shown below) is on quite small scales
Dessart & Owocki 2003, A&A, 406, L1
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No expectation of porosity from simulations
Natural explanation of line profiles without invoking porosity
Finally, to have porosity, you need clumping in the first place, and once you have clumping…you have
your factor ~3 reduction in the mass-loss rate
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No expectation of porosity from simulations
Natural explanation of line profiles without invoking porosity
Finally, to have porosity, you need clumping in the first place, and once you have clumping…you have
your factor ~3 reduction in the mass-loss rate
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No expectation of porosity from simulations
Natural explanation of line profiles without invoking porosity
Finally, to have porosity, you need clumping in the first place, and once you have clumping…you have
your factor ~3 reduction in the mass-loss rate (for Pup, anyway)
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f = 0.1
f = 0.2
f = 0.05
f ~ 0.1 is indicated by Ha, UV, radio
free-free analysis
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f = 0.1
f = 0.2
f = 0.05
And lack of evidence for porosity…leads us to
suggest visualization in
upper left is closest to
reality
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f = 0.1
f = 0.2
f = 0.05
Though simulations
suggest even smaller-scale
clumping