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 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:

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. 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: 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 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 a Cam (O9.5Ia) with wings extending to 1500 km/sec (Wilson) Expanding Shell WR Star Emission Absorption 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 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)) “It can be derived that” Here, a 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 Mass loss rate Terminal velocity Velocity law Geometry Time variability Described by simple scaling laws: Reimers mass loss formula

AGB Mass Loss 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 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 a 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 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