shirley by xiangpeng

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									 A Brief Summary of Star
Formation in the Milky Way



          Yancy L. Shirley


     Star Formation Disucssion Group
          April 1 2003 (no joke!)
                     Outline
Brief overview of Milky Way Star Formation (SF)
  Where? How much? How long?
  Molecular cloud lifetime & support

Dense Cores = sites of SF
  Compare & Contrast low-mass vs. high-mass
  Dichotomy in understanding SF across mass spectrum
  IMF cores to stars

Observational Probes
  Molecules & dust

Future Disucssion Topics
              SF in the Milky Way
1011 stars in the Milky Way
  Evidence for SF throughout history of the galaxy (Gilmore 2001)
SF occurs in molecular gas
  Molecular cloud complexes: M < 107 Msun (Elmegreen 1986)
  Isolated Bok globules      M > 1 Msun (Bok & Reilly 1947)
SF traces spiral structure (Schweizer 1976)

                                                          M51 Central
                                                            Region




                                    NASA
 SF Occurs throughout the Galaxy
Total molecular gas = 1 – 3 x 109 Msun (CO surveys)
  SF occurring within central 1 kpc
  SF occurring in outer galaxy > 15 kpc (Combes 1991)
  SF occurring nearby
     Rho Oph 120 pc, Lupus 130 pc, Taurus 140 pc, Orion 400 pc
     Pleiades 70 pc
SF occurs in isolated & clustered modes
          BHR-71                             W42




                                         Blum, Conti, &
    VLT                                  Damineli 2000
         Molecular Cloud Lifetime
Survey of CO towards clusters
   Leisawitz, Bash, & Thaddeus 1989
   All cluster with t < 5 x 106 yrs have molecular clouds M > 104 Msun
   Clusters older than t > 107 yrs have molecular clouds M < 103 Msun
   Lower limit to molecular cloud lifetime


Some young clusters show evidence for SF over periods of
t > 108 yrs (Stauffer 1980)

Lifetimes of 107 to 108 yrs
       Molecular Cloud Structure
Molecular clouds structure complicated:
  Clumpy and filamentary
  Self-similar over a wide range of size scales (fractal?)
  May contain dense cores: with n > 106 cm-3
  Transient coherent structures?
                                                      Lupus
        Serpens
                     Optical Av       Optical Av




                           L. Cambresy 1999
                                Gravity
Jeans Mass
   Minimum mass to overcome thermal pressure (Jeans 1927)
                         3/ 2
             kT 
 M Jeans  
            m G 
                                1/ 2  18M sunT 3 / 2 n 1/ 2
               H 


Free-fall time for collapse
                  1/ 2
         3 
 t ff  
         32G 
                         3.4 107 n 1/ 2 yrs
              
   n = 102 cm-3 => free-fall time = 3 x 106 yrs
   n = 106 cm-3 => free-fall time = 3 x 104 yrs
    Jeans Mass
0.5 1 2
          5
          10
            20
              50
               100
                 200
                   500
                     1000
            Star Formation Rate
Current SFR is 3 +/- 1 Msun yr -1 (Scalo 1986)

Assuming 100% SF efficiency & free-fall collapse
  Predicted SFR > 130 – 400 Msun yr -1 (Zuckerman & Palmer
  1974)
  TOO LARGE by 2 orders of magnitude!


SF is NOT 100% efficient
  Efficiency is 1 – 2% for large molecular clouds


All clouds do not collapse at free-fall
  Additional support against gravity: rotation, magnetic fields,
  turbulence
                SFR per unit Mass
  Assume LFIR ~ SFR, then SFR per unit mass does not
vary over 4 orders of magnitude in mass (Evans 1991)
     Plot for dense cores traced by CS J=5-4 shows same lack of
   correlation (Shirley et al. 2003)
     Implies feedback & self-regulation of SFR ?
              Rotational Support
Not important on large scale (i.e., molecular cloud
support)
   Arquilla & Goldsmith (1986) systematic study of dark clouds
   implies rotational support rare

Rotational support becomes important on small scales
   Conservation of angular momentum during collapse
      Results in angular momentum problem & solution via
      molecular outflows
   Spherical symmetry breaking for dense cores
      Formation of disks
   Centrifugal radius (Rotational support = Gravitational support)
   (Shu, Admas, & Lizano 1987) :

                            G 3 M 3 2
                       Rc 
                              16a 8
                 Magnetic Support
Magnetic field has a pressure (B2/8) that can provide
support
   Define magnetic equivalent to Jeans Mass (Shu, Adams, & Lizano 1987):
    M cr  0.13G 1/ 2  B dA  103 M sun B / 30G R / 2 pc
                                                                  2



   Equivalently: Av < 4 mag (B/30 mG) cloud may be supported
   M > Mcr “Magnetically supercritical”
   Equation of hydrostatic equilibrium => support perpendicular to B-field
                                  1  2        
                                             
        2
     d r              1
     2  P                    2  B  ( B  ) B
     dt               0                                
Dissipation through ambipolar-diffusion increases timescale for
collapse (Mckee et al. 1993):
              3
    t AD            7.3  10 13 xe yrs
           4G ni
   Typical xe ~ 10-7 => tAD ~ 7 x 106 yrs
Observed Magnetic Fields




                           Crutcher 1999
                     Turbulent Support
Both rotation & magnetic fields can only support a cloud in
one direction

Turbulence characterized as a pressure:
   Pturb ~ vturb2

General picture is turbulence injected on large scales with a power
spectrum of P(k) ~ k-a
   Potentially fast decay t ~ L / vturb => need to replenish

Doppler linewidth is very narrow:            2 ln 2kT               T
                                    Dv  2             0.22km / s
                                                 m                 mamu
   CO at 10K Dv = 0.13 km/s
   Low-mass regions typically have narrow linewidth => turbulence decays
   before SF proceeds?
   High-mass regions have very large linewidths
       CS J=5-4 <Dv> = 5.6 km/s
Rho Oph Dense Cores




              Motte, Andre, & Neri 1998
  Low-mass Dense Cores
                                        B335
N 2H + J = 1 - 0




                                        10,000 AU




                                    IRAS03282




            Caselli et al. 2002   Shirley et al. 2000
Star Formation within Cores
Orion Dense Cores
CO J=2-1




  VST, IOA U Tokyo
                     Lis, et al. 1998
Dust Continuum: Dense Cores
  350 m                         350 m




           Mueller et al. 2002
          High-mass Dense Cores                  RCW 38

   M8E                 S158         Optical




   W44                S76E          Near-IR




CS J = 5-4, Shirley et al. 2003   J. ALves & C. Lada 2003
       High-mass: Extreme Complexity

S106




                  Near- IR
                  Subaru
                                       H2
Orion-KL Winds & Outlfows
             SF in Dense Cores
Star formation occurs within dense molecular cores
  High density gas in dense cores (n > 106 cm-3)
  Clumpy/filamentary structures within molecular cloud
     SF NOT evenly distributed
  Low-mass star formation may occur in isolation or in clustered
  environments
     Low-mass defined as M_core < few Msun
  High-mass star formation always appears to occur in a clustered
  environment
Average Properties:
  Low-mass: R < 0.1 pc, narrow linewidths (~ few 0.1 km/s)
  High-mass: R ~ few 0.1 pc, wide linewidths (~ few km/s)

There is a dichotomy in our understanding of low-mass
and high-mass protostar formation and evolution
Low-mass Evolutionary Scheme




                       P.Andre 2002
 Low-mass: Pre-protostellar Cores
Dense cores with no known internal luminosity source
   SEDs peak longer than 100 m
  Study the initial conditions of low-mass SF
     B68                                        L1544
                            SCUBA 850 m                ISO 200 m




                                  10,000 AU


                                                   Ward-Thompson et al. 2002

                               3.5’ x 3.5’                 12’ x 12’
        High-Mass Star Formation
Basic formation mechanism debated:
   Accretion (McKee & Tan 2002)
      How do you form a star with M > 10 Msun before radiation pressure
      stops accretion?
   Coalescence (Bonnell et al. 1998)
      Requires high stellar density: n > 104 stars pc-3
      Predicts high binary fraction among high-mass stars


Observational complications:
   Farther away than low-mass regions = low resolution
   Dense cores may be forming cluster of stars = SED dominated by
   most massive star = SED classification confused!
   Very broad linewidths consistent with turbulent gas


Potential evolutionary indicators from presence of :
   H2O, CH3OH masers
   Hot core or Hyper-compact HII or UCHII regions
High-mass Evolutionary Sequence ?




            A. Boonman thesis 2003
        UCHII Regions & Hot Cores
 UCHII Regions and Hot Cores observed in some high-
mass regions such as W49A
          VLA 7mm Cont.                       BIMA




        DePree et al. 1997               Wilner et al. 1999
Chemical Tracers of Evolution?
  High Mass Pre-protocluster Core?

Have yet to identify initial
configuration of high-mass star
forming core!
   No unbiased surveys for such
   an object made yet


Based on dense gas surveys,
what would a 4500 Msun, cold
core (T ~ 10K) look like?

Does this phase exist?




           Evans et al. 2002
           IMF: From Cores to Stars
    dN/dM ~ M-1.6 – 1.7 for molecular clouds & large CO
  clumps
     dN/dM ~ M-2.35 for Salpeter IMF of stars
     How do we make the stellar IMF ?

  Rho Oph (60 clumps):
dN/dM ~ M-2.5, M>0.8 Msun
(Motte et al. 1998)


 Serpens:
dN/dM ~ M-2.1 (Testi &
Seargent 1998)
  CO: Molecular Cloud Tracer
 Hubble                        CO J=3-2
Telescope                      Emission




                                          CSO
  NASA, Hubble Heritage Team
   Dense Gas Tracers: CS & HCN
                                              CS 5-4


CO 1-0        CS 2-1           HCN 1-0




         Helfer & Blitz 1997             Shirley et al. 2003
 Comparison of Molecular Tracers
Observations of the low-mass PPC, L1517 (Bergin et al.)
Astrochemistry




 E. F. van Dishoeck 2003
       Dust Extinction Mapping

Good pencil beam probe for Av up to 30 mag (Alves et al
1999)
           Dust Continuum Emission

Optically thin at long
wavelengths => good
probe of density and
temperature structure

     ~ 1 at 1.2 mm for
    Av = 4 x 104 mag


    Dust opacities
    uncertain to order of
    magnitude!



 SCUBA map of Orion
Johnstone & Bally 1999
                Some Puzzles
              Based on question in Evans 1991

How do molecular clouds form?
  Does the same process induce star formation?
What is the relative importance of spontaneous and
stimulated processes in the formation of stars of various
mass?
What governs the SFR in a molecular cloud?
What determined the IMF evolution from molecular cloud
clumps to stars?
Do stars form in a process of fragmentation of an overall
collapse?
Or rather, do individual stars form from condensed
regions within globally stable clouds?
                More Puzzles
How do you form a 100 Msun star?
Is high-mass SF accretion dominated or coalescence
dominated?
  Does the mechanism depend on mass?
What are the initial conditions for high-mass cluster
formation?
How does SF feedback disrupt/regulate star formation?
  Outflows, winds, Supernovae
What is a reasonable evolutionary sequence for high-
mass star forming regions?
IS SF in isolated globules spontaneous or stimulated?
Are we actually observing collapse in dense core
envelopes?

								
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