1. 1 Lecture 2: Gas Cooling and Star Formation

2. How to make a galaxy Create Hydrogen, Helium and dark matter in a Big Bang Allow quantum fluctuations to cause some regions to be denser than others Ensure a large amount of dark matter, so there is enough mass to ensure the dense regions collapse due to gravity Add Dark Energy so the Universe expands at correct rate THESE ARE THE BASIC INGREDIENTS OF COSMOLOGY THE PILLARS ARE: Abundances of Light Elements The Cosmic Microwave Background The Large Scale Structure in the Distribution of Galaxies The Expanding Universe

3. How to make a galaxy Create Hydrogen, Helium and dark matter in a Big Bang Allow quantum fluctuations to cause some regions to be denser than others Ensure a large amount of dark matter, so there is enough mass to ensure the dense regions collapse due to gravity Gas cools to central regions of dark matter halos Dens gas then forms stars Add Dark Energy so the Universe expands at correct rate

4. Dark matter particles are collisionless and only detectable in bulk by their gravitational influence Baryons radiate - a sure sign that dissipative processes are at work. Dissipative means baryonic matter can loose energy by radiative processes, resulting in a loss of thermal energy from the system. So the baryonic component shrinks within the dark matter halos. Once the gas is stabilised by thermal pressure, loss of energy by radiation is an effective way of decreasing internal pressure, allowing region to contract and re-establish pressure equilibrium. Gas cools, settles into a disk, surrounded by a dark matter halo Dissipation

5. Important Cooling Processes Type Reaction Name Free-Freee – + X + → γ + e – + X + Bremsstrahlung Free-Bounde – + X + → X + γRecombination Bound-Freee – + X → 2e – + X + Collisional Ionization Bound-Bounde – + X → e – + XCollisional Excitation Electron-Photonγ + e – → γ + e – Inverse Compton Scattering

6. When relativistic electron collides with a low energy photon, e.g. the CMB. Imparts energy to the photon. The gas therefore cools. Inverse of Compton scattering where photon imparts energy to a slow electron. Important when T electrons >> T photons Important at: High redshifts CMB: T ~ 2.73(1+z)K γ + e – → γ + e – As the Universe expands, cooling times become long, so unimportant at low redshifts. X-rays from inverse Compton Scattering are also commonly seen in supernovae and active galactic nuclei (AGNs).

7. Bremsstrahlung Happens in very hot gasses where the atoms already ionised. A charged particle is accelerated around the nucleus of an ionised atom. The strong electromagnetic attraction alters the course of the charged particle. Important at Temperatures: Primordial* Gas > 10 6 K Enriched** Gas > 10 7 K *primorial gas: hydrogen + helium **enriched gas: additional metals from stars/supernovae e – + X + → γ + e – + X +

8. 1. Collisional Excitation Atoms are excited by collisions with electrons, then radiatively decay to the ground state At lower Temperatures (but still > 10 4 K) several processes occur. Each is Temperature dependent as well as metallicity dependent e – + X → e – + X

9. 2. Collisional Ionization Atoms ionized by collisions with electrons: kinetic energy equal to ionisation threshold is removed form the gas. Gas therefore cools. e – + X → 2e – + X + 3. Recombination Atoms ionized by collisions with electrons: kinetic energy equal to onisation threshold is removed form the gas. Gas therefore cools. e – + X + → X + γ

10. Fine Structure If gas is enriched, collisions with neutral hydrogen and the few free electrons excite fine structure levels of low ions such as OI, CII. At low Temperatures (< 10 4 K) cooling rates drop as most electrons have recombined. Neutral Hydrogen If gas is primordial, cooling only proceeds (slowly) via formation of H 2 which proceeds via gas phase reactions such as: e – + H 0 → γ + H − followed by H − + H 0 → H 2 + e – and/or H + + H 0 → H 2 + + γ then H 2 + + H 0 → H 2 + H +

11. Cooling rates can be calculated: Primordial Gas

12. Cooling rates can be calculated: Enriched Gas

13. Cooling rates can be calculated: More detailed determination Assumes collisional ionization equilibrium Assumes photo-ionization equilibrium

14. Mass of Collapsed Halos Using Press-Schechter or simulations, we know the mass of halos that have collapsed at various redshifts. We know the virial temperatures (see table). So we can use our knowledge of cooling times to determine whether gas will cool to the centre of the halos. To collapse we require that the cooling time is less than the free fall, or dynamical time, t cool < t dyn In this case energy can be removed sufficiently quickly to allow rapid gravitational collapse to take place.

15. Cooling Time Dynamical Time For a sphere: ~ 6.5 × 10 9 f ½ n ⅓ yr where f is the gas fraction of the cloud and n is the density. This uses crude approximations to the cooling rates from the various processes Λ is the cooling rate

16. Cooling Time versus Dynamical Time Thus setting t cool = t dyn we get: T 6 3/2 ~ f -1/2 n -1/2 ≈ 2.5 Using the Virial theorem to relate Temperature & density to the Mass of the dark matter halos, we find that: M = 1.2 ×10 13 T 6 3/2 f 3/2 n -1/2 (units of solar masses) Thus, according to our cooling rates, dissipational collapse will occur when M ≤ 3 ×10 13 f 2 For small f this gives us approximately Galaxy Masses

17. Virial Temperature and Pressure Support Gas experiencing a strong virial shock has its kinetic infall energy thermalized and is heated to the virial temperature Pressure supports the gas against gravitational collapse. But radiative cooling processes allow gas to dissipate energy Conservation of Angular Momentum results in gas forming a disk

18. Angular Momentum in Dark Matter Halos We can Calculate the distribution of angular momentum of dark matter halos Bullock et al 2001

19. Angular Momentum in Disks

20. Successes of simple model! Model disk formation using or simple assumptions of cooling of gas that was shock heated to the virial temperature Simple assumption: the disk scale-lengths are related to the radial size and angular momentum of the dark matter halos : r dis k λ r virial e.g. Successfully reproduce a population of galaxies that match the observed Tully-Fisher relation Relates the rotational velocity of galaxies (Vc) to their Magnitude (M)

21. Cold Flows How good are our assumptions? Does gas really shock at the virial radius??

22. How good are our assumptions? Does gas really shock at the virial radius?? Certainly, structure formation is not spherical in CDM

23. Gas Virial Radius (extent of Dark Matter) Gas in simulations does not shock Cools along the filamants ``Cold Mode” Accretion Cold Flows

24. Brooks et al. 2009 Cold gas can stream all the way into the disk region without ever shock heating Cold Mode accretion dominates for galaxies with mass < 2.5×10 11 M * See Keres et al 2005 ~

25. Cold Mode Accretion

26. Stars form in centre of Dark Matter Halos Star Formation

27. Gas Fragments into Giant Molecular Clouds (GMCs) where star formation occurs

28. Collapse of GMCs A cloud, becomes unstable and begins to collapse when it lacks sufficient gaseous pressure support to balance gravity (Jeans Crtieria) and when shear forces in the disc are not too large (Toomre Criteria).

29. The Jeans Mass A cloud, becomes unstable and begins to collapse when it lacks sufficient gaseous pressure support to balance gravity. When the sound-crossing time is less than the free-fall time, pressure forces win and the system bounces back to equilibrium. However, when the fee-fall time is less than the sound crossing time, gravity wins and the region collapses.

30. The Jeans Mass Assume GMCs are self gravitating, homogeneous, isothermal sphere: Free-fall tiime Sound crossing time For

31. The Jeans Mass Then the Jeans Mass is R J is ½ the Jeans length

32. The Toomre Criteria Large regions will be torn apart by the shear faster than they tend too collapse, namely faster than the gravitational free-fall time. These regions within disks are able to collapse if: Q = csκcsκ πGΣ < 1 Where κ is the epicycle frequency κ = √2 ( Vc2Vc2 R2R2 VcVc R dV c dR ) + ½ and Σ is the surface density of the gas

33. Star Formation Since gas is required for star formation, it is logical to look at the relation between the Star Formation rate (SFR) and the surface density of gas Star formation in spiral galaxies have shown the Schmidt law to be a surprisingly good description

34. Star Formation

38. The Initial Mass Function (IMF) N(M) ∝ M -α dM (e.g. Kroupa 2001) N(M) ∝ M -2.35 dM for M>M sun (Salpeter 1953) Bulk of mass integral Consider an ensemble of stars born in a molecular cloud (single stellar population) The distribution of their masses can be described piecewise by power-laws