1. Lecture 15 main sequence evolution

2. Recall: Initial cloud collapse A collapsing molecular cloud starts off simply:  In free-fall, assuming the pressure gradients are too small to have much effect  The gas is approximately isothermal, if gas is optically thin so energy can be efficiently radiated away. The time it takes for the shell containing mass M r to collapse to r=0 is the free- fall time scale:

3. Fragmentation Clouds don’t collapse to form one giant star. Stars tend to form in groups or clusters As isothermal collapse progresses, the increase in density means the Jeans mass must decrease. Thus if the cloud has some density inhomogeneities, smaller clumps within the cloud may collapse later on.

4. Adiabatic cloud collapse We have assumed an isothermal collapse, which is approximately true if energy can be radiated away efficiently. As the density increases, however, the gas becomes optically thick. At the other extreme, no energy is transported out of the cloud, and the collapse is adiabatic. In this case, the temperature rises as the cloud collapses. Thus the Jeans mass increases during collapse. The point where fragmentation stops depends on the point when the collapse becomes adiabatic.

5. Minimum Jeans mass Make a crude estimate of the lower mass limit of fragmentation.

6. Other effects Our estimate is crude, but gives an idea of the basic principles involved. We have neglected:  In using the Jeans criterion, we were assuming a static cloud, and not accounting for the energy of collapse  The details of radiative transport, vaporization of dust grains, dissociation of molecules, ionization of atoms.  Rotation: this will lead to the formation of a disk  Departure from spherical symmetry  Magnetic fields A fully numerical calculation (first done by Richard Larson, 1969) can improve the precision by including these and other physical effects.

7. Evolution of a solar mass protostar 1.Isothermal collapse begins 2.Collapse becomes adiabatic, and cloud heats up.

8. Evolution of a solar mass protostar 1.Isothermal collapse begins 2.Collapse becomes adiabatic, and cloud heats up. 3.A quasi-stable protostar forms, radiating gravitational potential energy 4.Collapsing material continues to accrete onto a disk 5.Core collapse continues when dust vapourizes and H 2 dissociates.

9. Evolution of a solar mass protostar 1.Isothermal collapse begins 2.Collapse becomes adiabatic, and cloud heats up. 3.A quasi-stable protostar forms, radiating gravitational potential energy 4.Collapsing material continues to accrete onto a disk 5.Core collapse continues when dust vapourizes and H 2 dissociates. 6.Nuclear fusion begins, and an energetic jet expels material and angular momentum from the system

10. Herbig-Haro objects Jets associated with star formation interact with the surrounding ISM, exciting the gas and forming bright, emission line objects. These are HH objects. Here we can actually see the stellar disk, illuminated by the central, obscured, star A movie of a HH jet

11. Stellar disks Young main sequence stars often still have disks, even after the molecular cloud has been dispersed. Infrared-emitting dust disk around  -Pic. The central star has been subtracted. The dust disk around Vega. At least one large planet is known to exist within this disk.

12. Quasi-stable collapse The quasi-static evolution of a protostar depends on the timescale for liberating gravitational potential energy, compared with the free-fall timescale. Recall the Kelvin-Helmholtz timescale: If the collapsing protostar is heated to ~3000 degrees, where do you expect it to lie on the HR diagram, relative to its future main sequence position?

13. The Hayashi track The Hayashi track represents a boundary set by convection.  Cooler stars are unstable because the energy cannot be transported rapidly enough.

14. ZAMS More massive stars are hot enough to begin burning Carbon  Leave Hayashi track at higher luminosity  Evolve very quickly; will be born and die (explode) before solar mass stars even make it to the ZAMS

15. Zero age main sequence The point where stars first reach the main sequence and begin equilibrium hydrogen burning is the Zero Age Main Sequence (ZAMS). The time taken to reach the ZAMS is inversely related to mass Mass (M Sun )Time to reach ZAMS (10 6 years) 15.00.0617 9.00.1505 5.00.5759 3.02.514 2.255.855 1.518.21 1.2529.45 1.050.16 0.5155.0

16. Break

17. Recall: main sequence structure The three divisions in a stellar interior are the nuclear burning core, convective zone and radiative zone. Energy, in the form of gamma-rays, is generated solely in the nuclear burning core. Energy is transferred towards the surface either in a radiative manner or convection depending on which is more efficient at the temperatures, densities and opacities.

18. The end of the MS: Solar mass stars The main sequence phase ends with the depletion of hydrogen in the core.  Isothermal He core  hydrogen burning continues in a shell around the He core.

19. Solar mass stars The H-shell burning phase produces even more energy than the core-burning phase  R and L increase, but T decreases  This is the subgiant branch When the isothermal core gets massive enough, it cannot support the envelope and collapses (on the Kelvin-Helmholtz timescale)

20. Degenerate cores When the density of a gas gets sufficiently high, the Pauli exclusion principle applies: no two fermions (e.g. electrons) can occupy the same quantum state.  Pressure becomes due to the nonthermal motions of electrons, and is independent of temperature A solar mass star’s core will collapse when it reaches ~13% of the total stellar mass. Lower mass stars have more degenerate cores: thus the cores may never collapse before the next stage of nuclear burning begins.

21. More massive stars Cores are convective, so composition is homogeneous When the hydrogen fraction decreases enough, the entire star contracts, releasing some gravitational potential energy.  The luminosity increases slightly  Temperature increases (since radius decreases).  This marks the end of main sequence evolution.

22. Eddington luminosity Very massive stars are unstable, due both to core pulsations and due to their very high luminosities. Thus there is a maximum luminosity at which it is possible to remain in hydrostatic equilibrium. This is the Eddington luminosity: If the luminosity exceeds this limit, mass loss must occur.

23. Stellar evolution Stars on the main sequence have a finite lifetime Red giants, red supergiants, and white dwarfs are all examples of stars that have evolved off the main sequence. The luminosity and effective temperature are the main observable predictions of stellar models.

24. Red Giants During the first few million years after leaving the main sequence, the changes in the core cause a rapid expansion of the envelope, leading to a much larger, and therefore cooler, star.

25. Stellar evolution: overview As the core loses pressure, it collapses and heats up(~2 million years) New energy sources cause envelope to expand and luminosity to increase, producing red giants (~20 million years)

26. Stellar evolution: overview Red giants lose most of their envelopes in an expanding shell of low density gas (planetary nebula, ~10,000 years)

27. Stellar evolution: overview The remnant core cools and fades to become a white dwarf (which lives for billions of years).