Transcript L and T Dwarfs

Mass Loss and Winds
• All massive stars with L>104LSun have winds
• All low and intermediate stars in post AGB
evolution (M>0.6 MSun or L>103.7LSun have
winds
• Lower mass, less luminous stars do too, but
mass loss rates are lower
• Evidence for winds includes
– P Cygni profiles
– Radio or IR excesses
– Thermal x-ray emission
The Solar Wind
•
The only cool dwarf with a directly observable wind, BUT
– We do see x-rays from solar type stars (hot gas similar to the Sun’s
corona (T~106, ~108 cm-3)
– Observed decrease in rotation with age implies torque from a wind
– ISM around nearby stars is disturbed (by wind…)
•
Solar wind discovered from space observations in the early ’60s
– Wind inferred earlier from aurora, geomagnetic storms, comet tails
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Parker (1958) demonstrated the wind must be there as a
consequence of the high temperature of the corona and the low
density of the ISM
Coronal gas pressure decreases outward, but tends toward a finite
pressure at infinite radius (and that finite pressure is larger than
the ambient pressure in the ISM)
Wind can be described by conservation of mass and momentum
combined with the ideal gas law:
dv 2a 2  (GM * r )

dr
v2  a2
a  (p /  )1 2  soundspeed
If the Star Rotates…
• Early work of Mufson’s (Mufson & Liszt 1975) describes the
effect of rotation on an expanding stellar wind assuming the
corona rotates uniformly at a given angular velocity and
angular momentum is conserved.
r dv r

vr dr
2a 2  v  (GM * / r )
2
vr  a 2
2
• Rotation increases the speed of the wind due to centrifugal
force
• In the case of the Sun, virtually all of the driving force of
the wind comes from the pressure gradient
• Using appropriate values for the solar corona, we get wind
speeds of about the right size for the solar wind.
• Can study the evolution of the solar wind by varying the
rotational velocity, and using solar evolutionary models to
estimate solar radius and temperature, and phenomenological
model of magnetic field
Evolution of the Solar Wind
• Initially high mass loss rate decreases as rotation,
magnetic field decline
• Increases again in the old Sun due to enlarging radius
Figure from
MacGregor, ESO
Workshop, 1997
Three Solar Winds
• Wind from open magnetic fields
– Terminal velocity= 700-800 km/sec
– nion=2.7 cm-3
– 1.5 x 10-14 Msun per year
• Wind from closed magnetic fields
– Terminal velocity = 400 km/sec
– nion=7 cm-3
– 2 x 10-14 Msun per year
• Coronal mass ejections
– Terminal velocity = ~1000 km/sec
– nion=0.2 cm-3
– Uncertain mass loss rate
Hot Star Winds
•
“Hot” means hydrogen
is ionized.
Hot star winds are important from
several astrophysical contexts:
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Evolution of massive stars
Physical state of the ISM
Chemical evolution of galaxies
Interpretation of integrated spectra
of extragalactic star bursts
– Phenomenon itself is interesting
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Early work (1940’s)on spectroscopy
of WR stars – mass loss by high
speed outflows, stratification of
ionization (Beals)
Winds understood in the context of
P Cygni profiles
Winds in O stars also inferred from
C III l5696A emission in  Cam
(O9.5Ia) with wings extending to
1500 km/sec (Wilson)
Expanding Shell
WR Star
Emission
Emission
Winds in the Far UV
• Winds became accessible to study with the beginning of far UV
astronomy – discovery of P Cygni profiles for C IV, Si IV and N
V in Orion supergiants
• Optical lines appear only in the case of extreme winds, but UV
lines are formed in the weak-wind approximation (modelling the
circumstellar continuum as a geometrically diluted continuum of
a static photosphere)
• Decoupling radiative transfer from equations of statistical
equilibrium and equations of gas dynamics allowed progress
– Winds in OB supergiants are spherically symmetric, terminal
velocities greater than stars’ escape velocities
– Single-fluid approximations are OK (no differential streaming)
– NOT analogs of the solar wind – existence of ions in the flow means
temperature not too hot. Winds dynamically “cold,” driven by
radiation pressure
– Non-radiative heating also necessary
• Modern treatments advanced to strong-wind limit with full
physics included
Modeling Hot Star Winds
• Radiation driven wind theory developed by Lucy
and Solomon (1970) and Castor, Abbot, and Klein
(CAK, 1975) to explain FUV line profiles
• Unified Model Atmospheres for hot stars spherically extended, sub- and supersonic
atmospheric structure, including mass-loss rate,
density, and velocity structure
• Calculate energy distributions and spectra
simultaneously for photospheric and wind lines,
and mixed cases
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Fully line-blanketed models
millions of spectra lines
equilibrium solutions with 150 ionic species
hydrodynamic equations for radiation-driven winds
• Detailed spectrum synthesis in UV and optical
Observed vs. Synthetic Spectrum
of LMC Sk –68 137 (O3)
Flux Distributions
• Free-free thermal emission causes IR
and radio excesses (x10)
• Bound-free edges in UV and EUV also
affected
– Change in density stratification
– Velocity shifts in resonance lines
weakens optical thickness, more UV flux
• More ionizing flux in the ISM
Wind Momentum-Luminosity Relation
• Hot star winds result from radiation pressure
• Properties of stellar winds must reflect the
luminosities of stars
• Wind momentum (mass loss rate x terminal
velocity) should be a function of the photon
momentum rate (f(L/c))
.
M v 
• “It can be derived that”
1 1
L
0.5
R*
• Here,  is the power law of the line strength
distribution function
• Wind momentum depends on luminosity and weakly
on radius
• Wind momentum-luminosity relation can be used as
a distance indicator
Areas of Current Work
• Variability of hot star winds
• Time scales of less than an hour to
several days
• Rotational modulation – wind
structures tied to photospheric origin
• Pulsations
• Non-spherical winds
Winds in Cool Stars
• Observed as blue-shifted cores in the
center of optically-thick resonance lines,
emission features in UV resonance lines
• Basic physical parameters include
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Mass loss rate
Terminal velocity
Velocity law
Geometry
Time variability
• Described by simple scaling laws: Reimers
mass loss formula
.
 ( L LSun )
13
M  4 x10
M Sun / yr
( g / g Sun )(R / RSun )
AGB Mass Loss
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The “lemming” diagram
Derived from Bowen
models
Argues that Reimers
formula is an artifact
of a selection effect
We observe “Reimers”
mass loss because
stars with higher rates
of mass loss are very
short-lived
Wind Properties in Cool Stars
• Geometry – we know winds aren’t spherically symmetric
• Shocks, discontinuities, non-monotonic flows, circulation
patterns
• Magnetic properties – as yet no field measurements
• Wind acceleration mechanisms not yet known
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Convection driven acoustic waves
Pulsation driven acoustic waves
Alfven waves
Radiation pressure on dust
Parker-type thermal pressure gradients
Steady vs. stochastic?
• NLTE in the ionization equilibria, need to infer mass loss
rates (e.g Mg II emission strength depends on both the
mass loss rate and the ionization state of Mg)
• Turbulence, turbulent pressure gradients
• Binary systems
Wind in  Tra
• Hybrid-chromosphere
star
• High temperature
emission lines
• High velocity winds (P
Cygni profile in Mg II
• Fit with
• spherical flow
• Chromosphere at base
of wind
• Simple wind
parameterization
Complexities in Winds
• Recall the 3 solar winds
• Wind properties may depend on
– Source of wind
– Corotating interaction regions (when a fast wind overtakes a
slower wind)
– Properties may depend on viewing angle
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Models also too simplistic
NLTE
Not in hydrodynamic equilibrium
More complex radiative transfer needed
Gas not in radiative equilibrium (heating and cooling times
longer than dynamical time scales
• Must solve radiative transfer and dynamics simultaneously