# Presentation: 085820

## Star Formation at Very Low Metallicities

- Jonathan Tan
- Dept. of Astronomy
- University of Florida

- Primordial Mini-Halo (Abel ea.) 			        Orion Nebula Cluster (McCaughrean)

## Importance of First Stars and their Mass

- Metal Enrichment
- proto-galaxies and IGM
- Illumination
- SN & GRBs?
- Progenitors of SMBHs?

- CMB polarization (WMAP Page et al. 06)
- H 21cm (LOFAR Morales & Hewitt 03)

- Z of low Z halo stars (e.g. Christlieb et al. 03)
- Z of Lya forest?(Schaye et al. 03 Norman et al. 04)

- NIR bkg. intensity (Santos et al. 03;
- 		Fernandez & Komatsu 06)
- NIR bkg. fluctuations (Kashlinsky et al. 04)

- Observations

- JWST	(Weinmann & Lilly 2005)
- SWIFT	(Bromm & Loeb 2002)

- Sloan	(Fan et al. 03; Willott et al. 03)

- Influence on structure formation: SMBHs, Globular Clusters, Galaxies? Critical Z ~ 5x10-3Z (Bromm et al. 2001)

- “Simple” problem: initial conditions, chemistry,
- no feedback from other stars, weak B-fields(?)

## Star Formation: A Complicated, Non-linear Process

- Computers cannot yet solve the complete problem: time step *t*  1/(*G*)1/2
- Need to break it up into separable parts, tractable with analytic and simple numerical methods (e.g. Stahler, Shu, Taam 1980). Need to understand physics, make reasonable approximations, compare to observations and numerical simulations.

- Physics: gravity vs pressure (thermal, magnetic, turbulence, radiation, cosmic rays).
- Equation of state: heating and cooling, decay and sources of turbulence, diffusion of B-fields, generation of B-fields (dynamo), etc.
- Chemical evolution of dust and gas.

- Wide range of scales (~10 dex in space, time) and multidimensional.
- Uncertain initial conditions/boundary conditions.

- Numerical models

-   Analytic theory

-    Observations

- PopIII Initial Conditions

## Outline

- The formation of zero metallicity stars
- Initial conditions
- Accretion and protostellar evolution
- Feedback
- Implications for the next generation
- Ionization feedback and HD
- Critical metallicity
- Transition to present-day star formation

## Numerical Simulations: Results

- Abel, Bryan, Norman (2002)

- Form pre-galactic halo ~106M
- Form quasi-hydrostatic gas core inside halo:
- 	M≈4000M,  r ≈10 pc,
- 	nH ≈10 cm-3,  fH2 ≈10-3, T~>200K
- Rapid 3-body H2 formation at nH>~1010cm-3.
- 	Strong cooling -> supersonic inflow.

- 4.  1D simulations (Omukai & Nishi 1998): Form quasi-hydrostatic protostar nH ≈1016-17cm-3, T ≈2000 K: optically thick, adiabatic contraction -> hydrostatic core with m*≈0.005M, r*≈14R (also Ripamonti et al. 2002)

- Numerical Simulations: Results

## The initial conditions for primordial star formation

- H2 cooling ->
- Tmin ~ 200 K,   ncrit ~ 104 cm-3
- MBE = 380M   cs=1.2 km/s

- Centrally concentrated cloud, inefficient cooling -> quasi-hydrostatic contraction
- Density structure: ~self-similar,
- r-2.2

- Yoshida et al. 2006

## Slow contraction damps turbulent motions and inhibits fragmentation

*z* = 19

- *t* = 9 Myr

- *t* = 0.3 Myr

- *t* = 30 kyr

- *t* = 3 kyr

- *t* = 1.5 kyr

- *t *= 200 yr

## Heating by Dark Matter Annihilation?

- Abel, Bryan, Norman (2002)

- baryons

- dark
- matter

- Chen & Kamionkowski 2004; Ripamonti ea. 2007; Spolyar ea. 2007; A.Natarajan & Tan 2007

- Main uncertainties: DM(r<10-3pc), m, <v>, ftrap
- Could potentially delay/halt collapse

## The Accretion Rate:Collapse of an Isentropic Sphere

- Density structure: self-similar,              k≈2.2
- ~singular polytropic sphere in virial and hydrostatic equilibrium
- Can model with a polytropic equation of state
-                                    		        =1.1

## Comparison withNumerical Models

- Omukai & Nishi (1998)

- Tan & McKee (2004)

- Ripamonti et al. (2002)

- ABN (2002)

- Note: ABN curve based on estimate from infall rate *prior* to protostar formation

- Within factor 2 of Bromm & Loeb 03 for 1st 104 yr  (M<50 M)

- Yoshida et al. (2006)

- Comparison withNumerical Models

## Collapse to a disk around a central protostar

- Rotation: core forms from mergers
- and collapse along filaments:
- expect J>0
- fKep  vcirc / vKep  ~ 0.5

## Disk Models

- Surface density

- Thickness

- Ionization

- Temperature
- Tc , Teff

- Toomre

- =0.3

- m*

- .

- 17x10-3M/yr

- 6.4x10-3M/yr

- 2.4x10-3M/yr

- (Tan & McKee 2004; Tan & Blackman 2004)

## Magnetic Fields: Inverse Helical Dynamo

- Turbulence in disk amplifies field strength, eventually saturating due to back reaction of B-field (Tan & Blackman 2004).
- Large-scale helicity is generated in each hemisphere, leading to strong, coherent, large-scale B-fields, which rise up to form a corona and outflow.

- Blackman & Field (2002)

## Magnetic Fields: Saturation Points

- Exponential growth of seed field (B~10-16G) eventually becomes resistively limited (Blackman & Field 2002).
- t ~ m*/m*

- .

- Unsaturated

- Saturated

- Growth limited by local accretion time in disk
- (ss=0.01, 0.3)

- Growth limited by disk lifetime

- Tan & Blackman (2004)

- Gravito-Turbulence

- MRI-Turbulence

- Either Gravito- or MRI-driven turbulence can provide viscosity to remove angular momentum.

## Evolution of the Protostar depends on Accretion Rate

- follow energy of
- protostar as it
- accretes

- Deuterium burning for Tc>106K
- Structural rearrangement after tKelvin
- Eddington model for 
- Solve for r*(m*), until reach main sequence

- Assume polytropic stellar structure and continuous sequence of equilibria

- (Stahler et al. 1980; Palla & Stahler 1991; Nakano 1995; Omukai & Palla 2003; Tan & McKee 2003)

## Evolution of the Protostar

- Initial condition
- m* = 0.04 M
- r* = 14 R
- (Ripamonti et al. 02)

- Protostar is large (~100 R) until it is older than tKelvin
- Contraction to Main Sequence
- Accretion along Main Sequence

- Comparison with Stahler et al. (1986), Omukai & Palla (2001)

- : Radius

## Evolution of the Protostar

- : Luminosity

- Boundary Layer

- Accretion Disk

## Evolution of the Protostar

- :Ionizing
- Luminosity

- fKep=0.5

- Spectrum depends on initial rotation
- c.f. Omukai & Palla (2003)

## Overview of Radiative Feedback

## Growth of the HII Region with Rotation

- Balance ionizing flux vs
- recombinations and infall

- Find stellar mass at breakout
- rHII = rg, where vesc=ci
- polar; equatorial

- HII Region

- McKee & Tan 2007

## Disk Photoevaporation

- Mass loss rate:
- for zero age main sequence

- .

- Hollenbach et al. 94

- thin neutral disk

## Accretion rate reduced by expansion of HII region

- Accretion continues from regions shadowed by the accretion disk. Need to calculate vertical structure (i.e. thickness) of disk.

## Radial profile of disk 

- m*=100M

## Feedback-limited accretion

- Self-consistent model for
- growth and evolution of
- protostar, including:
-  declining accretion rate
-  accretion disk (r,z)
-  protostellar structure
-  ionizing feedback

| K’ | Ti,4 | m*f
(M) |
| --- | --- | --- |
| 1 | 2.5 | 164 |
| 0.5 | 2.5 | 77 |
| 2.0 | 2.5 | 365 |

- Expect masses ~60-400M depending on K’, fKep

- McKee & Tan 2007

## Implications of IMF for supernovae and metal enrichment

- However, rapid rotation can lead to significant mass loss (Meynet et al. 2004, 2006; Hirschi 2006), so initial stellar mass is not necessarily the final pre-SN mass.

- Fate of NON-ROTATING metal-free stars
- (Heger & Woosley 2002)

## Impact of Massive PopIII Stars

- Radiative and Mechanical Feedback
- 	- Unbinding of gas from host and neighboring mini halos (ionization, winds, SNe)
- 	(Shapiro ea. 2004; Kitayama ea. 2004; Whalen ea. 2004; Iliev ea. 2005; Alvarez ea. 2006; Yoshida ea. 2007)
- 	- Dissociation of H2 while star shines -> contribution to FUV background
- 	(Haiman ea 1997, 2000; Omukai & Nishi 1999; Glover & Brand 2001; Machacek ea. 2001; Oh & Haiman 2002; Yoshida ea. 2003; Ahn & Shapiro 2006; Susa 2007; O’Shea & Norman 2007b)
-  	- (Relic) Ionization can promote H2 and HD formation; promote star formation
- 	(Shapiro & Kang 1987; Kang & Shapiro 1992; Susa ea 1998; Uehara & Inutsuka 2000; Nakamura & Umemura 2002; Ricotti ea. 2002; Nagakura & Omukai 2005; Johnson & Bromm 2006, 2007; Yoshida 2006; Ripamonti 2006; Susa & Umemura 2006; Yoshida ea. 2007a,b)
- 	-Cosmic ray ionization can promote H2 and HD formation
- 	(Stacy & Bromm 2007; Jasche ea. 2007)

## Impact of Massive PopIII Stars

- 2.	Production of metals and dust to change gas cooling properties: m*f? SNe yields (m*f)? Dust formation in SNe? Mixing?
- 	- Supernova remnants
- 	(Madau ea. 2001; Mori ea. 2002; Bromm ea. 2003; Wada & Venkatesan 2003; Norman ea. 2004; Greif ea. 2007)
- 	- Effect on cooling function
- 	(Omukai 2000; Bromm ea. 2001; Schneider ea. 2002; Bromm & Loeb 2003; Frebel ea. 2007; Smith & Sigurdsson 2007; Smith ea. 2007)

## Comparison to Local Massive Star Formation

- e.g. Beuther et al. (2007, PPV)

- Orion
- Nebula
- Cluster
- (McCaughrean)

## Overview of Physical Scales

- AV=7.5
- NH=1.6x1022cm-2
- =180 M pc-2

- AV=1.4
- NH=3.0x1021cm-2
- =34 M pc-2

## Overview of Physical Scales

- AV=7.5
- NH=1.6x1022cm-2
- =180 M pc-2

- AV=1.4
- NH=3.0x1021cm-2
- =34 M pc-2

- AV=200
- NH=4.2x1023cm-2
- =4800 M pc-2

## Overview of Physical Scales

- AV=7.5
- NH=1.6x1022cm-2
- =180 M pc-2

- AV=1.4
- NH=3.0x1021cm-2
- =34 M pc-2

- PopIII core mostly supported by thermal pressure. Weak cooling. No fragmentation.

- AV=200
- NH=4.2x1023cm-2
- =4800 M pc-2

## Transition from PopIII

- PopIII cores supported by thermal pressure
- Present-day cores and protoclusters supported by magnetic fields and turbulence
- Magnetic fields and turbulence unlikely to be effective in post-PopIII regions
- 	(B-fields probably too weak; turbulence damps out)

-   Post-PopIII regions may be able to cool
-     effectively via HD and/or metals-dust
-   If cs << virial velocity of halo, then gas
-     may collapse to form a rotationally
-     supported thin disk: possible fragmentation.
-     (Lodato & P. Natarajan 2006; Clark, Glover & Klessen 2007)
-      However, continued infall and merging of halos
-     may maintain turbulence (Wise & Abel 2007)

## Conclusions

- Approximately convergent initial conditions for star formation: set by H2 cooling. Mostly thermal pressure support + slow cooling
- -> no (disk) fragmentation -> single ~massive star in each mini-halo (no clusters of primordial stars). DM heating?

- Need analytic model to follow the growth of the protostar: accretion rate; accretion disk; protostellar evolution; feedback.
- Large (~AU) protostar, contracts to main sequence for m*>30M
- This is when feedback processes start to become important.

- Feedback processes depend on core rotation and accretion rate:
- Gradual reduction in SF efficiency because of ionizing and radiation
- pressure feedback for m*>30M. Final mass in fiducial case (K’=1, fKep=0.5), is ~160M, set by ionizing feedback on core and disk.

- Changing accretion rate by factor of ~5 (K’=0.5-2) leads to final masses in the range ~60 - 400M   (seeds for SMBHs?)

- Impact of PopIII stars on subsequent SF: suppressed cooling due to FUV background; enhanced cooling from metals/dust and ionization catalyzed H2/HD -> rotational supported, fragmenting cores? Eventual build up of B-field support to resemble present-day SF clusters.