Stellar winds and mass loss
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Stellar Winds and Mass Loss. Brian Baptista. Summary. Observations of mass loss Mass loss parameters for different types of stars Winds colliding with the ISM Effects on stellar evolution. Some History. Nova like objects are discovered Diagnostics of mass laws are generated for hot stars

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Stellar Winds and Mass Loss

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Stellar Winds and Mass Loss

Brian Baptista


  • Observations of mass loss

  • Mass loss parameters for different types of stars

  • Winds colliding with the ISM

  • Effects on stellar evolution

Some History

  • 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?

  • Two basic Parameters

    • The mass loss rate, , or the amount of mass loss per unit of time.

      • This is an important quantity for stellar evolution as stars with different mass loss rated will evolve differently.

    • The terminal velocity of the wind, , or the velocity the ejecta have at large distances from the star.

      • 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,


  • 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.

Terminal Velocity

  • 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,

Hey wait what is beta?

  • 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

  • P Cygni Profiles

  • Emission Lines

    • Ions

    • Molecules

  • Infrared and radio excesses

P Cygni Profiles

  • 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.)

  • 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”

  • Profiles caused by a strongly saturated line will give us the velocity profile of the region

    • Saturated lines are most sensitive to the velocity profile, because the Doppler core will give a hard edge at v∞

  • Profiles due to unsaturated lines can give us the mass loss rate

    • 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

  • The biggest advantage is that this can be used to study mass loss from the ground.

    • 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)

  • The lines typically have Doppler widths of a few hundred km/s

    • This is not the terminal velocity of the winds

    • These lines are formed near the star

  • The lines are typically formed by recombination

    • The emissivity is proportional to ρ2

    • These lines must be formed in regions of high density

Mass loss determination

  • 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

    • Any photon created by this process will escape the region

  • We can determine a total line luminosity

    • The mass loss will be determined by

Emission lines for Molecules

  • The same approach can be used for molecules around cool stars

    • 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

Infrared and Radio Excesses

  • Radio excess has only been measured for a few stars

    • As a 10-6M/yr would correspond to a few mJy

  • Infrared excesses have men more heavily observed

    • 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

  • O and B type stars

    • 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

Mass loss rates (continued)

  • Central stars in planetary nebulae have very low mass loss rates

    • Typical values are /yr and terminal velocities of 3000 km/s

  • Cool stars such as red super giants also have low mass loss rates

    • 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)

  • 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

  • Winds deposit enriched materials back into the ISM, and massive stars can even create dust particulate

  • Fast winds can collide with previously ejected winds

    • 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

    • Rotation and clumping can cause different shock structures in the ejecta

      • 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

  • Mass loss can cause changes in surface composition

    • 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

  • Formation of PNe

  • Lack of luminous red super giants

    • The massive stars loose so much mass that their have insufficient mass to become convective

  • Formation of white dwarves

    • Mass loss is responsible for stars that have masses less than 8M  not becoming SN, but instead becoming white dwarfs

    • The winds can remove up to 6.6 M  worth of material

Effects on evolution (continued)

  • 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

    • 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

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