Transcript Stellar Winds and Mass Loss

Stellar Winds and Mass Loss
Brian Baptista
Summary
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Observations of mass loss
Mass loss parameters for different
types of stars
Winds colliding with the ISM
Effects on stellar evolution
Some History
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Nova like objects are
discovered
Diagnostics of mass
laws are generated for
hot stars
Mass loss rated from
cool giants were
observed
Finally, time
dependant
mechanisms are
studied.
How do we define mass loss?
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Two basic Parameters
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The mass loss rate, M , or the amount of mass loss per
unit of time.
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The terminal velocity of the wind, v, or the velocity the
ejecta have at large distances from the star.
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This is an important quantity for stellar evolution as stars
with different mass loss rated will evolve differently.
Different ejection theories predict different velocities, so it
can be used to determine the ejection mechanism.
The energy deposited into the ISM per unit time is,
1  2
KE  M v
2
M-dot
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General form for a spherically
symmetric wind.
Use the mass continuity equation.
The same amount of gas per unit time
flows through a sphere at any
distance.
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M  4r  (r )v(r )
2
Terminal Velocity
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Gas that escapes from upper
atmospheres of stars, starts at small
radial velocity.
The gas is then accelerated to the
terminal velocity, at large r.
Often the terminal velocity is
approximated to,
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 ro ( R* , vo , v ,  ) 
 R* 
v(r )  vo  (v  vo )1    v 1 

r 
r
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Hey wait what is beta?
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Beta describes how
steep the velocity
profile is.
Hot stars have
steep profiles with
β=0.8
Cool stars have
smaller
accelerations β=2.0
Observations of Mass Loss
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P Cygni Profiles
Emission Lines
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Ions
Molecules
Infrared and radio excesses
P Cygni Profiles
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P Cygni is the prototype,
and was observed by Snow
and Morton in 1976
Most are observed using
UV resonance lines.
C IV, N V, and Si IV are use
in O to early B
C II is used in late B to A
Mg II is used in late B to M
P Cygni Profiles (cont.)
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The star emits a
continuum
The tube directly
between the star and
observer absorbs line
absorbs everything
between v=0-v∞
The region around the
star contains velocities
between -v∞ and v∞
“So what”
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Profiles caused by a strongly saturated line
will give us the velocity profile of the region
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Profiles due to unsaturated lines can give us
the mass loss rate
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Saturated lines are most sensitive to the velocity
profile, because the Doppler core will give a
hard edge at v∞
These profiles are fit using the above velocity
profile, with different numbers of absorbing ions,
until the profile matches the observed
unsaturated profile
The first group to make mass loss
determinations was Lamers and Morton in
1976
Emission Lines
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The biggest advantage is that this can
be used to study mass loss from the
ground.
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The star must have a high mass loss rate
on the order of 10-6M/yr
Most notable is Hα
Also, Paschen and Brackett lines of He II
Wolf-Rayet stars are dominated by lines
that form in high density winds
Emission Lines (continued)
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The lines typically have Doppler widths of a
few hundred km/s
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This is not the terminal velocity of the winds
These lines are formed near the star
The lines are typically formed by
recombination
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The emissivity is proportional to ρ2
These lines must be formed in regions of high
density
Mass loss determination
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Since the gas is expanding radial from the
star a photon that is created by
recombination will be created at a Doppler
shift that is greater than twice the thermal
width of the line
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Any photon created by this process will escape
the region
We can determine a total line luminosity
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The mass loss will be determined by
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M ~ v Ll
Emission lines for Molecules
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The same approach can be used for
molecules around cool stars
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The advantage is that they will from at large
distances from the star, 104R*.
CO J=1→0 lines are typically used
The velocities at this range are much lower than
the escape speed of the star, but they still
indicate mass loss
Knapp and Morris derived an expression for the
CO mass loss rate in 1985
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0.85
M ~ 5 1016 TB v 2 D 2 f CO
Infrared and Radio Excesses
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Radio excess has only been measured for a
few stars
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As a 10-6M/yr would correspond to a few mJy
Infrared excesses have men more heavily
observed
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IR emission is due to free-free emission within
1.5 stellar radii from the surface of the star.
These excess can be only a few tenths of
magnitudes
The mass loss rate from IR excess requires an
accurate determination of the velocity law
Mass loss rates
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O and B type stars
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These have been the most heavily studied
The terminal velocity of the ejecta is comparable
to the star’s escape velocity, but can depend on
the effective temperature of the star due to
radiation pressure
Krudritzki et al. determined that for galactic
stars, the loss rate is basically independent of
the stellar mass
 L* 
 
1/ 2 
log M v R*   1.37  2.07 log 6 
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 10 
Mass loss rates (continued)
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Central stars in planetary nebulae have very
low mass loss rates 
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Typical values are M ~10 M/yr and terminal
velocities of 3000 km/s
Cool stars such as red super giants also
have low mass loss rates
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6 G3 to M2 stars of class II and Ia that are in
binary systems have been measured
These are between 10-9 and 10-6 M/yr
The terminal velocities of 17 and 160 km/s
Mass loss rates (continued)
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AGB stars have extremely
high mass loss rates
The rate is linked to the
period of pulsation of the
stars
The rate seems to saturate
at about 10-4M /yr
These however have low
terminal velocities of 5-25
km/s
Interactions with the ISM
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Winds deposit enriched materials back into the ISM,
and massive stars can even create dust particulate
Fast winds can collide with previously ejected winds
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These can explain hot bubble around hot stars, ring
nebulae around WR stars, and ultra compact HII regions,
as well as PNe
The time evolution of different models can be used
to create a range of different out comes
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Rotation and clumping can cause different shock
structures in the ejecta
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Rotation can cause a higher density mass loss region in the
equatorial regions
Clumping can cause mass loading, and slow the shocks
down
Effects on evolution
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Mass loss can cause changes in surface composition
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Formation of PNe
Lack of luminous red super giants
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When the outer layers of an atmospheres are blown off, it
exposes the convective cores of the stars
These core will show an extreme over abundance of heavy
elements
The massive stars loose so much mass that their have
insufficient mass to become convective
Formation of white dwarves
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Mass loss is responsible for stars that have masses less than 8M
 not becoming SN, but instead becoming white dwarfs
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The winds can remove up to 6.6 M  worth of material
Effects on evolution
(continued)
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For stars with masses greater
than 30M mass loss can
change the amount of time
that a star spends on the main
sequence
Since the mass is so large for
these stars throughout the
main sequence lifetime
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The luminosity over the
lifetime of the star can
change
The lower luminosity means
that the star will have a
longer MS lifetime
The final mass that the star
will have is effected