THERMAL STABILITY OF LOW MASS STARS
Let us consider a star in a hydrostatic equilibrium, i.e. we have d2 r/dt2 = 0. Let the star be also in
a thermal equilibrium, i.e. we have dS/dt = 0. Now we shall make a perturbation that is so slow that
is does not di

Physics of Star and Planet Formation
Exercises #4 Due 30 April 2012 @ 09:45 pm
Name:_Date: _
Show your work on attached sheets. Write clearly and explain
your physical reasoning.
Attached you find the observed magnitudes/flux densities for 4 young sources

Physics of Star and Planet Formation
Exercises #6 Due 30 May 2012
Name:_Date: _
Show your work on attached sheets. Write clearly and explain
your physical reasoning.
Imagine NASA's Terrestrial Planet Finder has just flown and returned the following
images

Physics of Star and Planet Formation
Exercises #3 Due 9 April 2012 @ 09:45 am
Name:_Date: _
Show your work on attached sheets. Write clearly and explain
your physical reasoning.
1. What is the crossing time for a star cluster of radius 1 pc for a velocity

RED GIANTS
There is a large variety of stellar models which have a distinct core envelope structure. While
any main sequence star, or any white dwarf, may be well approximated with a single polytropic
model, the stars with the core envelope structure may

ACCRETION DISKS
A review article on accretion disks was published by J. E. Pringle (1981, Annual Review of
Astronomy and Astrophysics, 19 137). The classical papers on the structure of thin accretion disks
are: J. E. Pringle and M. J. Rees (1972, Astronom

MASS LUMINOSITY RELATION FOR MASSIVE STARS
Within the Eddington model Pg /P = const, and a star is an n = 3 polytrope. Large
mass stars have small , and hence are dominated by radiation pressure, and the opacity in them is
dominated by electron scattering

HIGH MASS STARS - EDDINGTON MODEL
The model was proposed by Eddington in the 1920s, when very little was known about physical
properties of matter in stellar interiors. The model assumes, that pressure is provided by perfect,
fully ionized gas and radiati

NUMERICAL INTEGRATIONS
In most cases of any interest the equations of stellar structure cannot be integrated analytically.
Instead, the integrations have to be performed numerically. With fast and inexpensive computers
this is and easy thing to do. You ma

OPACITY
A good description of opacity is provided by Schwarzschild in Structure and Evolution of Stars
(Chapter II) . In practical applications opacities calculated and tabulated by the Los Alamos group
are used (e.g. Cox, A. N., and Tabor, J. E. 1976, Ap

Radial Stellar Pulsations
Up to this point, you have studied stars in hydrostatic equilibrium. In the next few lectures, we
will consider oscillating stars. We begin with the simplest but perhaps most important case: radial
pulsators. These stars remain s

Astro 302/514
20
23 November 2004
HELIOSEISMOLOGY II: probing the solar interior
In the last lecture, we showed that the wavefunction for solar p and f modes has the form
P1
(r, , , t) = eit Ym (, ) (r ).
(1)
The frequency and radial wavefunction ) satisf