x-ray emission from massive stars david cohen dept. of physics & astronomy swarthmore college

X-ray Emission from Massive Stars

David CohenDept. of Physics & Astronomy

Swarthmore College


David CohenDept. of Physics & Astronomy

Swarthmore College

The work discussed here is with collaborators: Stan Owocki and Rich Townsend (U. Delaware), Asif ud-Doula (U. Delaware and Swarthmore), Maurice Leutenegger (Columbia), & Marc Gagne (West Chester)

Students: Roban Kramer (’03), Kevin Grizzard (St. John’s College, ’06), Casey Reed (’05), Stephanie Tonnesen (’03), Steve St. Vincent (’07)


O and early B stars: M > 8Msun; Teff > 20,000 K; term “massive stars” used interchangeably with “hot

stars”

OUTLINE

1. Introduction

a. Solar x-ray emission…vs. massive star x-ray emission

b. Massive stars and their winds

2. The wind-shock paradigm

3. Chandra spectroscopy of Puppis and Orionis: wind shocks

4. Chandra spectroscopy of 1 Orionis C: signatures of a magnetized wind

5. Conclusions

X-rays from the Sun

Remember - for thermal radiation - the frequency of light (the energy of each photon) is proportional to the temperature of the emitter:

Human body = 300 K 10 microns, or 100,000 Å (infrared)

Sun, light bulb filament = 6000 K 5000 Å (visible, yellow)

Hot star’s surface = 40,000 K 750 Å (far ultraviolet)

Really hot plasma = 5,000,000 K 6 Å (X-ray)

*don’t forget that thermal emitters give off photons with a range of wavelengths; those listed above represent the peak of the distribution or the characteristic wavelength.

Note: an Angstrom unit (Å) is equivalent to 0.1 nanometers (nm)

The Sun is a strong source of X-rays(10-5 of the total energy it emits)

It must have ~million degree plasma on it

The hot plasma is generally confined in magnetic structures above – but near - the surface of the Sun.

Visible solar spectrum: continuum, from surface

X-ray/EUV solar spectrum: emission lines from hot, thin

plasma above the surface

We can use spectroscopy - in our study of massive stars (where spatial structure can’t be imaged) - to diagnose plasma kinematics

(via Doppler-broadened line shapes) and plasma location with respect to the stellar surface (via UV-sensitive line ratios)

X-ray/EUV solar spectrum: emission lines from hot, thin

plasma above the surface

Theme: spectroscopy as a proxy for imaging.

This hot plasma is related to magnetic fields on the Sun: confinement, spatial structure, conduits of

energy flow, heating

More magnetic structures on the Sun:

x-ray image from TRACE

The Sun’s magnetic dynamo requires rotation + convection to regenerate and

amplify the magnetic field

Sunspots over several days: rotation

Note granulation, from convection

TRACE composite

OK, so the Sun emits x-rays - quite beautifully - and they’re associated with its magnetic

activity, related to convection and rotation…

But what of hot, massive stars?

OUTLINE

1. Introduction






5. Conclusions

Hot, Massive Stars

Stars range in (surface) temperature from about 3500 K to 50,000 K

Their temperatures correlate with mass and luminosity (massive stars are hot and very bright): a 50,000 K star has a million times the luminosity of the Sun (Tsun = 6000 K)

Stars hotter than about 8000 K do not have convective outer layers - no convection - no dynamo - no hot corona…

…no X-rays?

Our Sun is a somewhat wimpy star…

Puppis:

42,000 K vs. 6000 K

106 Lsun

50 Msun

Optical image of the constellation Orion

Note: many of the brightest stars are blue (i.e. hot, also massive)

In 1979 the Einstein Observatory made the surprising discovery that many O stars (the hottest, most massive stars) are strong X-

ray sources

Note: X-rays don’t penetrate the Earth’s atmosphere, so X-ray telescopes must be in space

Chandra X-ray image of the Orion star forming region1 Ori C: a 45,000 K

O-type star

So, we’ve got a good scientific mystery: how do massive stars make X-rays?

Could we have been wrong about the lack of a magnetic dynamo - might massive star X-rays be similar to solar X-rays?

Before we address this directly, we need to know about one very important property of massive stars (that might provide an alternate explanation for their X-rays)…

Massive stars have very strong radiation-driven stellar winds

Hubble Space Telescope image of Car; an extreme example of

a hot-star wind

What is a stellar wind?

It is the steady loss of mass from the surface of a star into interstellar space

The Sun has a wind (the “solar wind”) but the winds of hot stars can be a billion times as strong as the Sun’s

How do we know these hot-star winds exist?

Spectroscopy!

blue wavelength red

Absorption comes exclusively from region F - it’s all blue-sifted

You can read the terminal

velocity (in km/s) right off the blue edge of the absorption line

rest wavelength(s) – this N V line is a doublet

Why do hot star winds exist?

The solar wind is actually driven by the gas pressure of the hot corona

But hot-star winds are driven by radiation pressure

Remember, photons have momentum as well as energy:

pE

chc

h

And Newton tells us that a change in momentum is a force:

F mamdv

dtdp

dt

So, if matter (an atom) absorbs light (a photon) momentum is transferred to the matter

Light can force atoms to move!

re, the radius of an electron, giving a cross section, T (cm2)

The flux of light, F (ergs s-1 cm-2)

The rate at which momentum is absorbed by the electron

dp

dtFc

LT

4cR2

arad LT

4cR2

By replacing the cross section of a single electron with the opacity (cm2 g-1), the combined cross section of a gram of plasma, we get the acceleration due to radiation

For a (very luminous) hot star, this can compete with gravity…but note the 1/R2 dependence, if arad > agrav, a star would blow itself completely apart.

And free electron opacity, and the associated Thompson scattering, can be significantly augmented by absorption of photons in spectral lines – atoms act like a resonance chamber for electrons: a bound electron can be ‘driven’ much more efficiently by light than a free one can (i.e. it has a much larger cross section), but it can only be driven by light with a very specific frequency.

Radiation driving in spectral lines not only boosts the radiation force, it also solves the problem of the star potentially blowing itself apart:

As the radiation-driven material starts to move off the surface of the star, it is Doppler-shifted, making a previously narrow line broader, and increasing its ability to absorb light.

0

cont.

Optically thick line – from stationary plasma (left); moving plasma (right) broadens the line and

increases the overall opacity.

The Doppler desaturation of optically thick (opaque) lines allows a hot-star wind to bootstrap itself into existence!

And causes the radiation force to deviate from strictly 1/R2 behavior: the radiation force on lines can be less than gravity inside the star but more than gravity above the star’s surface.

OUTLINE

1. Introduction






5. Conclusions

X-rays from shock-heating in line-driven winds

The Doppler desaturation that’s so helpful in driving a flow via momentum transfer in spectral lines is inherently unstable

The line-driven instability (LDI) arises when a parcel of wind material is accelerated above the local flow speed, which moves it out of the “Doppler shadow” of the material behind it, exposing it to more photospheric radiation, and accelerating it further…

Numerical modeling of the hydrodynamics show lots of structure: turbulence, shock waves, collisions between “clouds”

This chaotic behavior is predicted to produce X-rays through shock-heating of some small fraction of the wind.

A snapshot at a single time from the same simulation. Note the discontinuities in velocity. These are shock fronts, compressing and

heating the wind, producing x-rays.

There are dense inter-shock regions, though, in which cold material provides a source of photoelectric absorption

Even in these instability shock models, most of the wind is cold and is a source of x-ray continuum opacity - x-rays emitted by the shock-heated gas can be absorbed by the

cold gas in the rest of the wind

Keep this in mind, because it will allow us to learn things about the physical properties of

a shocked wind via spectroscopy

X-ray line profiles can provide the most direct observational constraints on the

x-ray production mechanism in hot stars

Wind-shocks : broad lines

Magnetic dynamo : narrow lines

The Doppler effect will make the x-ray emission lines in the wind-shock scenario broad, compared to

the x-ray emission lines expected in the coronal/dynamo (solar-like) scenario

OUTLINE

1. Introduction






5. Conclusions

So, this wind-shock model - based on the line-force instability - is a plausible alternative to the

idea that hot star x-rays are produced by a magnetic dynamo

This basic conflict is easily resolved if we can measure the x-ray spectrum of a hot star at high

enough resolution…

In 1999 this became possible with the

launch of the Chandra X-ray Observatory

Now, for some data

Pup(O4 I)

10 Å 20 Å

N V

O VIIO VIII

Ne XSi XIV

Fe XVII

Ne IX

Mg XII

Focus in on a characteristic portion of the spectrum

Ne X Ne IX Fe XVII

Pup

(O4 I)

12 Å 15 Å

Capella - a cooler star: coronal/dynamo source

Differences in the line shapes become apparent when we look at a single line (here Ne X Ly)

Pup

(O4 I)

Capella (G2 III)

The x-ray emission lines in the hot star Pup are broad -- the wind shock scenario is looking good!

But note, the line isn’t just broad, it’s also blueshifted and asymmetric…

lab/rest wavelength

We can go beyond simply wind-shock vs. coronal:

We can use the line profile shapes to learn about the velocity distribution of the shock-heated gas and even its spatial distribution within the wind, as well as learning something about the amount of cold wind absorption (and thus the amount of cold wind).

What Line Profiles Can Tell Us

The wavelength of an emitted photon is proportional to the line-of-sight velocity:

Line shape maps emission at each velocity/wavelength interval

Continuum absorption by the cold stellar wind affects the line shape

Correlation between line-of-sight velocity and absorption optical depth will cause asymmetries in emission lines

The shapes of lines, if they’re broad, tells us about the distribution and velocity of the hot plasma in the wind -- maybe discriminate among specific

wind shock models/mechanisms

We will now build up a physical – but flexible – empirical x-ray emission line profile model:

Accounting for the kinematics of the emitting plasma (and the associated Doppler shifting/broadening);

Radiation transport (attenuation of the line photons via bound-free absorption in the cold wind component).

Note that our line-profile model, while physical, is agnostic regarding the heating mechanism.

Emission Profiles from a Spherically Symmetric, Expanding Medium

A uniform shell gives a

rectangular profile.

A spherically-symmetric, x-ray emitting wind can be built up

from a series of concentric shells.

Occultation by the star removes red photons, making

the profile asymmetric

Continuum Absorption Acts Like Occultation

Red photons are preferentially absorbed, making the line asymmetric: The peak is shifted to the

blue, and the red wing becomes much less steep.

wavelength

redblue

Contours of constant optical depth (observer is on the left)

The model has four parameters:

: v(r) (1 R /r)

Ro,q : j 2r q

: (p0;z) dz'

r'2 (11

r')z

for r>Ro

The line profile is calculated from:

Increasing Ro makes lines

broader; increasing * makes them more blueshifted and skewed.

Ro=1.5

Ro=3

Ro=10

=1,2,4

where

M

4Rv

L 8 2 je R

1

1 r2drd

Line profiles change in characteristic ways with * and Ro, becoming broader and more skewed with increasing * and broader and more flat-topped with increasing Ro.

A wide variety of wind-shock properties can be

modeled

Ro=1.5

Ro=3

Ro=10

=1,2,4

In addition to the wind-shock model,

our empirical line-profile model can also describe a corona

With most of the emission concentrated near the photosphere and with very little acceleration, the resulting line profiles are very narrow.

We fit all the unblended strong lines in the Chandra spectrum of Pup: all the fits are

statistically good

Ne X

12.13 ÅFe XVII

15.01 ÅFe XVII

16.78 Å

Fe XVII 17.05 Å

O VIII 18.97 Å

N VII 24.78 Å

We place uncertainties on the derived model parameters

Here we show the best-fit model to the O VIII line and two models that are marginally (at the 95% limit) consistent with the data; they

are the models with the highest and lowest * values possible.

lowest * best * highest *

Lines are well fit by our three parameter model: Pup’s x-ray lines are consistent with a spatially distributed, spherically symmetric, radially accelerating wind scenario, with reasonable parameters:

*~1 :4 to 15 times less than predicted

Ro~1.5

q~0

But, the level of wind absorption is significantly below what’s expected.

And, there’s no significant wavelength dependence of the optical depth (or any parameters).

The results for Pup were published several years ago, with Roban Kramer (Swarthmore 2003) as the lead author.

However, it’s generally been considered that other massive stars’ x-ray spectra were not consistent with the wind-shock scenario.

Much of the work shown on the next few slides – on Ori – was done by Kevin Grizzard (St. John’s College 2006)

Here’s another way of looking at the situation:

There are claims in the literature that the emission lines of most massive stars can be fit by Gaussian profiles.

We fit strong lines in the Chandra spectra of Ori with unshifted Gaussians (top), shifted Gaussians (center), and the wind profile model (bottom).

94%

73%

54%

Rejection probabilities are shown on the right of each panel.

Fit results for Ori summarized (to appear in Monthly Notices of the R.A.S., 2006)

*The wind profile model provides statistically good fits to all the lines. The onset radii (left) are exactly what’s expected from

the standard wind-shock picture. There is evidence for attenuation by the cold wind (right), but at levels nearly 10 times lower than expected. This is the same result that we

found for Pup.

Ro of several tenths to one stellar radius is expected based on numerical simulations of the line-force instability (self-excited on the left; sound

wave perturbations at the base of the wind on the right)

This is consistent with the results of the He-like f/i ratio analysis

So…what’s going on with the much lower wind optical depths?

The atomic cross sections are quite well known.

Could the mass-loss rates of massive stars be overestimated? By an order of magnitude? There would be very serious evolutionary implications (for, e.g., supernovae and chemical enrichment of galaxies).

There are, in fact, other recent papers that show several independent lines of evidence that wind mass-loss rates are lower than previously thought.

Some of these rely on the insight that clumping will cause density-squared diagnostics to overestimate mass-loss rates.

Density-squared processes – H-alpha emission (driven by recombination) and radio free-free emission – are commonly used to determine wind mass-loss rates.

Bouret, Lanz, & Hillier (2005): detailed fits to UV spectra;Puls et al. (2006): H-alpha and radio free-free excess;Fullerton, Massa, & Prinja (2006): FUSE and Copernicus P V absorption line fitting

12 X 1=1 12 X 1=1

12 X 1=1 12 X 1=1

emissivity, j = n2 X Vol

The effect of clumping on density-squared emission

j = 4

42 X 1=16

02 X 1=0

02 X 1=0

02 X 1=0

j = 16

Clumping’s effect on density-squared emission is scale-free (only the density contrast between clumps and the inter-clump medium matters).

However, we have begun to investigate a separate effect – porosity – the ability of photons to more easily escape through low-density inter-clump channels.

We have discovered that the key parameter for describing the reduction in effective opacity due to porosity is the ratio of the clump size scale to the volume filling factor.

We dub this quantity the porosity length, h.

Winds with porosity length increasing to the right

It turns out that line profiles (left) are not significantly affected until the porosity length is comparable to the stellar radius (unity, in the unitless formulation of these slides).

This degree of porosity is not expected from the line-driven instability. The clumping in 2-D simulations (right) is on quite small scales.

Note: these clumps are spherical.

The line-driven instability might be expected to compress clumps in the radial direction: pancakes, oriented parallel to the star’s surface. We’ve started working on models with

non-isotropic/oblate clumps: the Venetian-blind model.

There’s one more powerful x-ray spectral diagnostic that can provide useful information to test the wind-shock scenario:

Certain x-ray line ratios provide information about the location of the x-ray emitting plasma

Distance from the star via the line ratio’s sensitivity to the local UV radiation field

g.s. 1s2 1S

1s2s 3S1s2p 3P

1s2p 1P

resonance (r)

intercombination (i)forbidden (f)

10-20 eV

1-2 keV

Helium-like ions (e.g. O+6, Ne+8, Mg+10, Si+12, S+14) – schematic energy level diagram

The upper level of the forbidden line is very long lived – metastable (the transition is dipole-forbidden)

g.s. 1s2 1S

1s2s 3S1s2p 3P

1s2p 1P

resonance (r)


10-20 eV

1-2 keV

1s2s 3S

1s2p 3P1s2p 1P

resonance (r)


g.s. 1s2s 1S

While an electron is sitting in the metastable 3S level, an ultraviolet photon from the star’s photosphere can excite it to the 3P level – this decreases the intensity of the forbidden line

and increases the intensity of the intercombination line.

UV

1s2s 3S

1s2p 3P1s2p 1P

resonance (r)


g.s. 1s2s 1S

The f/i ratio is thus a diagnostic of the strength of the local UV radiation field.

UV

1s2s 3S

1s2p 3P1s2p 1P

resonance (r)


g.s. 1s2s 1S

If you know the UV intensity emitted from the star’s surface, it thus becomes a diagnostic of the distance that the x-ray

emitting plasma is from the star’s surface.

UV

Si XIII line complex in the Chandra spectrum of a massive star where the local UV mean intensity is not strong enough to affect

the forbidden-to-intercombination ratio.

r i f

r i f

Si XIII line complex in the Chandra spectrum of a massive star where the local UV mean intensity is strong enough to affect the

forbidden-to-intercombination ratio.

Here the f/i ratio is reduced, due the effects of UV photoexcitation… this occurs because the x-ray emitting plasma is relatively close to the photosphere.

We have fit line profile models simultaneously to the f-i-r complexes in four hot stars – and get consistent fits: Hot plasma smoothly distributed throughout the wind, above roughly 1.5 Rstar –

•The f/i line ratios are consistent with this spatial distribution

•The line profile shapes are also consistent with this distribution (as already was shown for single, unblended lines)

Conclusions for most massive stars: normal O-type supergiants

Spherically symmetric, standard wind-shock scenario describes the Chandra data for Pup and Ori (and, it looks like, most other massive stars too) – x-ray line profiles and line ratios are consistent with the expected distribution of hot plasma

There’s evidence for attenuation by the cold wind component, but the level of continuum absorption in the wind must be reduced from expected values by factors of ~10 (mass-loss rate reduction? some porosity?)

x-ray emission from massive stars david cohen dept. of physics & astronomy swarthmore college

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