IV. Star Formation
1. Basic Aspects of Star Formation • Star Formation: The fundamental cosmic (baryonic) process Determines cosmic fate of normal matter
Galaxy formation, Star Conditions for life evolution, IMF Formation
Elements Planet formation (He => U)
Clusters Light, K.E. of ISM
black holes WS 2018/19 (AGN,CDE stellar) IV 2
1 Star formation visible on local scales
3 Stellar associations: Does the energy feedback self-regulate the SF?
2. Star Formation in Spiral Galaxies
IC5333 in Hα star formation follows the gaseous spiral arms. 4
2 Star formation can also happen in diffuse, less dense HI gas at the rim of gas disks.
NGC 4625
UV disk is 4x larger than optical disk. Gil de Paz et al. (2005)
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3 2.1. “Kennicutt/Schmidt law”: 2 global relat.s
n SF threshold at g 6 M SF g (1989) ApJ, ˙
n 1.4
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4 1.40.15 (2.5 0.7) 104 g M yr 1kpc 2 SF 2 s 1Ms pc
g SF 0.017 g g dyn
WS 2018/19 CDE IV 9 Kennicutt (1998) ApJ, 498
Can we understand the SF – gas surface density relation?
if large - scale SF is produced by small - scale dynamics over a large area : g g 1.5 (for const. H) SF 1 g SF 2 (G g )
but : SF could be very short WS 2018/19 CDE IV 10
5 Starbursting galaxies drop off the KS correlation; the same for mergers.
Hensler (2012) ApSS Proceed. + WS 2018/19 CDE IV Kuehtreiber (2010) BSc. thesis11
2.2. Toomre‘s criterion for large-scale SF c Q 1.4 Gravitational instability of 3.36 G integrated vertical disk leads to SF
0.7 c Above a critical density the disk is crit 3.36 G gravitationally unstable.
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6 2.3. SF dependence on dense cloud cores
Kennicutt‘s law is valid for the total amount of gas:
1.4 SF HI
But SFR correl.s strongly with the dense molecular gas mass
(from HCN) Gao & Solomon (2004) ApJ, 606 WS 2018/19 CDE IV 13
Genzel et al. (2010)
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8 3. Giant Molecular Clouds
• Typical characteristics of GMCs: 4 6 – Mass = 10 ...10 M – Distance to nearest GMC = 140 pc (Taurus) – Typical size = 5..100 pc – Size on the sky of near GMCs = 5..20 x full moon – Average temperature (in cold parts)= 20... 30 K – Typical density = 103... 106 mol./cm3 – Typical (estimated) life time = ~107 year – Star formation efficiency = ~1%..10%
• Composition of material: – 99% gas, 1% solid sub-micron particles (‘dust’) (by mass) -4 -5 – Gas: 0.9 H2/H, 0.1 He, 10 CO, 10 other molecules (by number) – Dust: Mostly silicates + carbonaceous (< m in size) • Properties of the gas:
– Gas mostly in molecular form: hydrogen in H2, carbon in CO, oxygen in O (O2?), nitrogen in N2(?). – At the edges of molecular clouds: transition to atomic species. “Photo-Dissociation Regions” (PDRs).
– H2 cannot be easily observed. Therefore CO often used as tracer.
9 Nearby well-studied GMCs:
4 • Taurus (dist ≈ 140 pc, size ≈ 30 pc, mass ≈10 M): Only low mass stars (~105), quiet slow star formation, mostly isolated star formation. 4 • Ophiuchus (dist ≈ 140 pc, size ≈ 6 pc, mass ≈ 10 M): Low mass stars (~78), strongly clustered in western core (stellar density 50 stars/pc), high star formation efficiency 6 • Orion (dist ≈ 400 pc, size ≈ 60 pc, mass ≈ 10 M): Cluster of O-stars at center, strongly ionized GMC, O-stars strongly affect the low-mass star formation • Chamaeleon... • Serpens...
3.1. GMC mass distribution
Fukui & Kawamura (2010) Ann.Rev.A&A, 48, 547
10 3.2. typical GMCs in nearby galaxies
Fukui & Kawamura (2010) Ann.Rev.A&A, 48, 547 SMC GMCs LMC GMCs , with HI
M33 GMCs , with HI
IC10 GMCs with HI
Fukui & Kawamura (2010) Ann.Rev.A&A, 48, 547
11 3.3. Typical parameters of Galactic GMCs
parameter average range in
5 4 6 mass 5·10 5·10 … 5·10 M⊙ radius 20 10 … 50 pc
3 3 density 300 100 … 10 H2 /cm temperature 10 5 … 30 K sound speed 0.2 0.15 … 0.35 km/s turb. vel. dispersion 6 2 … 10 km/s rotation 0.3 ? km/s/pc
magn. field strength 50 20 … 100 Gauss lifetime 107 ? yr
free fall time 2·106 106 … 3·106 yr star formation efficiency 0.01 ?
9 PS. total amount of gas in GMCs in the Galaxy: M(H2) = 2·10 M⊙
3.4. Hierarchical Structure • Clump picture: hierarchical structure – Clouds (≥ 10 pc) – Clumps (~1 pc) • Precursors of stellar clusters – Cores (~0.1 pc) • High density regions which form individual stars or binaries
• Fractal picture: clouds are scale-free
V AD / 2 D 1.4 fractal dimension
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3.5. Molecular Cloud Structure
• Formation of IS clouds, • Filamentary structure, • Embedded dense clumps, • Turbulent dynamics, • Star Formation in densest clumps
McKee & Ostriker (2007) ARAA, 45
13 3.6. Molecular Cores
Most clumps don’t form stars. But if they do, they form many. Core mass spectrum is more interesting for predicting the stellar masses of the newborn stars.
Deep 1.3 mm continuum map of Ophiuchi (140 pc) at 0.01 pc (=2000 AU) resolution.
(Motte et al. 1998)
Clumps
14 3.7. Clump mass spectrum Orion B: First GMC systematically surveyed for dense gas and embedded YSOs by E. Lada 1990
Survey of gas clumps
Clumps in range M = 8..500 M
dN M 1.6 dM
dN MdN M 0.6 M 0.4 Most of mass in dln M dln M massive clumps
Core mass spectrum Result of survey: dN dN M 0.6 M (1.11.5) dln M dln M
for M < 0.5 M for M > 0.5 M
Motte et al. 1998
15 For a video see: http://www.astro.umd.edu/~ostriker/research/clouds/project.html
Larson 1981 3.8. Mol.Cloud Parameters
The larger and more massive the (sub- )clouds the larger the velocity dispersion (~T).
16 Fukui & Kawamura (2010) Ann.Rev.A&A, 48, 547
Larson 1981
17 2 Isolated gas and dense particle systems 3 n G M c E pot achieve Virial equilibrium: 2T + V = 0 5 2n Rc T: kinetic (+ thermal) energy M E c 2 V: potential energy th 2 assumptions: spherical cloud, 2G M 5 2n cloud mass M c c 2 L 3 n ρ(r) r-n 5/3, n 0 1.0, n 2
Larson 1981
Fukui & Kawamura (2010) Ann.Rev.A&A, 48, 547
18 The molecular gas fraction depends on the ISM pressure:
0.92 H2/HI ∝ P
after: Blitz & Rosolowsky, 2006, ApJ, 650
19 4. star formation in molecular clouds why not in atomic clouds? molecular clouds can cool to low temperatures (T = 10-30 K) and condense to high densities (n = 103-105 cm-3) atomic clouds cannot (T = 100-5000 K, n = 1-10 cm-3) low gas temperatures and high gas densities promote the possibility of star formation (Jeans criterium) some important molecules:
- diffuse clouds: CH, CN (1937, 1940) 12 13 18 - mol. clouds: H2, CO ( CO, CO, C O) + + - dense cores: H2CO, NH3, CS, HCN, HCO , H2D - maser sources: OH, H2O, SiO, CH3OH (methanol) - photodiss. regions: PAHs (10% of all carbon) dual role of dust: shielding molecular clouds from UV radiation + efficient cooling agent!!!
star formation in molecular clouds (ctd)
12 -4 - formation of H2: on dust surface CO/H2 ≈ 10 - destruction of H2: dissoc. energy 11.2 eV (optically thick) (binding energy 4.5 eV) 13 -6 - formation of CO: OH + C, CH + O, etc. CO/H2 ≈ 10 - destruction of CO: dissoc. energy 11.1 eV (optically thin) formation of giant molecular clouds (AV > 1): large-scale grav. instabilities (tidal limit) agglomeration of smaller clouds (spiral arms) turbulent compression of supersonic interstellar medium sources of turbulence: winds, supernovae, rotational shear destruction of molecular clouds:
photodiss. regions (PDR): dissoc. of H2 and CO HII regions: ionisation of HI
20 SF clouds Stars are formed in den dense cores of molecular clouds. Necessarily, they must be shielded against the IS radiation field, in particular, against UV. critical cloud radius
R τ 1 Z T 4 1013 dyn cm-3 cl 5 o pc Z 80 K P ext cloud cores cool (molecular lines) thermal instability Jeans unstable WS 2018/19 CDE IV 41 protostellar collapse
4. Star-formation Process
Hierarchical (fractal) structure of giant molecular clouds (GMCs)
protostellar fragments small scale condensations (cores) large scale condensations (clumps) diffuse molecular envelope warm atomic „skin“ H2
HI
cf. Surdin & Lamzin 1998, p. 99
21 Star Formation
Shrink size by 107; increase density by x 1021 ! Where planets also form
• Giant Molecular Cloud Core Raw material for star birth • Gravitational Collapse & Fragmentation Proto-stars, proto-binaries, proto-clusters
• Rotation & Magnetic Fields Accretion disks, jets, & outflows • Planets WS 2018/19Most may form in clusters!CDE IV C. Lada 43
4.1. Star formation by gravitational instability
simple approach
„collapse“
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22 4.2. Evolutionary sequence of the mol. clouds. The left panels are examples of Large Magellanic Cloud (LMC) giant molecular cloud (GMC) Type I (GMC 225, LMC N J0547-7014 in Fukui et al. 2008), Type II (GMC 135, LMC N J0525-6609), Type III (the northern part of GMC 197, LMC N J0540-7008) from the top panel.
Each panel presents Hα images from Kim et al. (1999) with GMCs identified by CO contours obtained with NANTEN: The contour levels are from 1.2 K km s−1 with 1.2 K km s−1 intervals. Open blue circles indicate the position of young clusters (Bica et al. 1996). The right panels are cartoon illustrations for each evolutionary stage. Open blue circles and filled red circles represent young clusters and HII regions, respectively.
4.3. Star-forming sites: The Orion GMC
the massive star factory
Orion Nebula From: CfA Harvard, Millimeter Wave Group (part of Orion GMC)
23 4.4. Bright IR sources
24 RCW 49 (Spitzer's infrared array camera)
Star-formation happens produces Star clusters.
Credit: NASA/JPL-Caltech, University of Wisconsin
4.5. Spectral SF indicators
UV of M31 presents young massive Optical image of M31 52 stars.
25 Inner rings with arms emerging from the ring. 54
26 Star Formation in the Whirlpool Galaxy
M83 lies in the southern constellation Hydra at a distance of bout 5 Mpc. (a) This visible-light image clearly shows the spiral arms. The presence of young stars and H II regions indicates that star formation takes place in spiral arms. (b) 21 cm. Note that essentially the same pattern of spiral arms is traced out in this image as in the visible-light photograph. (c) M83 has a much smoother appearance in this near-infrared view. This shows that cooler stars, which emit strongly in the infrared, are spread more uniformly across the galaxy’s disk. Note the elongated bar shape of the central bulge.
27 5. The SF – gas vol. density relation
k if SF is produced SF g 1) by growing perturbations : g g 1.5 SF 1 g SF 2 (G g )
1.5 2) by grav. collapse the same g 3) by cloud - cloud collisions : 2 SF cl
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SF consumtion and efficiency
by the gas consumption due to star formation dM dM SFR : * g M dt dt g
a characteristic gas consumption timescale cons can be expressed, as e.g. for M ( ) cons g cons dM 1: dt g M 0 M g ,0 g
M g (t) M g,0 exp(t cons )
The SF efficiency is then defined by :
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Facts The total energy of ionizing photons from massive stars over their lifetimes are comparable to the supernova type II energy. Star formation is visible as HII regions. Star formation visible as HII regions.
HII regions visible in Hα Lyc flux from massive stars: with IMF SFR Massive stars in UV: with IMF SFR Calibration of IR emission by dust SFR
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6. The Initial Mass Function
Salpeter (1955): γ = –2.35
Kroupa IMF:
log M/M WS 2018/19 CDE IV 66 Kroupa (2002)
32 One can derive the
Initial Mass Function dN(M) Mγ dM
Definition: (M) = dN(M)/dM Mγ = M-α normalized to
mu (M) dM =1 ml dN/d(logM) = dN/dM · M MΓ Γ=1+γ
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Similar to stellar IMF (Initial Mass Function)
Stellar IMF: Meyer et al. PP IV
Salpeter (1955) IMF: dN M 1.35 dln M
33 6.1. Cloud mass spectrum vs. IMF
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34 6.2. Massive Star Clusters
Massive Star Clusters
35 6.3. Gas infall triggers Starburst The case od NGC 1569: 6 • HI clouds (each ~10 M) fall towards in from a disk • 2 huge super star clusters are formed. (2005, AJ, 130)
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A counter-rotating HI envelope around NGC 4449 is dynamical- ly perturbed by the passage of DDO 125 and is condensed(?). Metallicity of infalling Gas?
WS 2018/19 CDE IV 75 Hunter et al. (1998)
36 An excess of massive stars in the local 30 Doradus starburst Spectroscopy of 247 stars with
>15 M in 30 Dor (upto 200 M). 32±12% more stars above 30 M
Schneider et al. 2018, Science, 359
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6.4. Low-mass SF sites
Brightness and velocity-integrated contours reflect hierarchical clumpy structure of molecular clouds visible in different species!
37 Initial Mass Function of Trapezium (Muench et al. 2001)
The massive-star HBL DBL
range is incomplete!
Completeness limit Completeness
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Tinker Bell Triplett: The “Bird” with extremely low SFRs in TDGs
Star formation in the tidal-tail -4 -3 blobs with rates ~ 10 … 10 M/yr WS 2018/19 CDE IV 80
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Star formation in the RPS blobs The detection of -4 -3 with rates of ~ 10 … 10 M/yr ESO 137-001
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m 5 More extreme case: MV ~ -9.4 , M* ~ 5.7 10 M SFR ~ 4.5 10-5 M /yr !!! Leo P
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39 7. Derivations of star-formation rates from spectroscopic SF indicators
without dust: −42 SFRH(M/yr) = 7.9 × 10 LHα(ergs/s) 41 = LH(ergs/s)/1.26 x 10 (Kennicutt, 1998) −42 = 5.5 × 10 LHα(ergs/s) (Calzetti et al., 2007) The variation of the calibration constant is ~ 15% for variations in Te = 5000-20000 K , 6 -3 and < 1% for variations in ne = 100-10 cm (Osterbrock & Ferland 2006).
-28 -1 -1 SFRFUV(M/yr) = 1.4 × 10 L FUV(erg s Hz )
with dust:
−38 0.885 40 44 SFR24μm(M/yr) = 1.31×10 (L24) (local, 1·10 < L24μm/ergs/s < 3·10 ) uncertainty of 0.02 in the exponent, 15% in the calibration constant
−43 42 43 = 2.04 × 10 L24 (global, 4·10 < L24μm/ergs/s < 5·10 ) (Calzetti et al., 2007) WS 2018/19 −42 CDE IV 83 SFR(M/yr) = 5.3 × 10 [LH + 0.031 L24μm]
Facts The total energy of ionizing photons from massive stars over their lifetimes are comparable to the supernova type II energy. Star formation is visible as HII regions. Star formation visible as HII regions.
HII regions visible in Hα Lyc flux from massive stars: with IMF SFR Massive stars in UV: with IMF SFR Calibration of IR emission by dust SFR Star formation is self-regulated. Massive stars clear-up their birthplaces. How to disentangle low SFRs
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40 7.1. Star SFRs derived from indicators (massive stars normalized to IMF) H and UV begin formation at -2 -1 to deviate below ~ 10 M yr . low rates Explanation: H preferably from higher- mass stars than UV IMF not complete in uppermost mass range.
WS 2018/19 CDE IV 85 Lee et al., 2009, ApJ, 706
Boselli et al. (G.H.)
H is only a necessary, but not a sufficient condition for SF!
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41 Can the IMF be global?
What are the consequences of low star-formation rates for the evolution of dwarf galaxies?
How is it treatable in numerical models of galaxy evolution?
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7.2. Possibilities to fill the IMF according to the SFR/cloud mass filled IMF reduced to star fraction IMF truncated at upper mass interval with N =1 *
star fractions!! Consequences of low SFR: Filled IMF: star fractions lead to SNII fractions heating Truncated IMF: longer lifetimes of heaviest stars; w/o SNeII? WS 2018/19 CDE IV 88
42 At low SFR 3 possibilities emerge: dN( m ) IMF : ( m ) ~ m dm
mu • a filled IMF can lead to N( m) N( m ) A ( m ) dm A ( m ) dm becoming fractions of 1 only! i.e. for massive stars m ml also N (m) mu SNII SFR : M ( m ) A m 1 dm * ml • The IMF is truncated condition : N( mu ) 1
• A stochatic IMF allows for individual massive stars
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MV of brightest star cluster vs. column SFR in various galaxies
Maximum star- cluster V brightness is correlated with the K-S SFR.
Exceptions are galaxies with starbursts, forming super star clusters.
Larsen, 2002, AJ, 124
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43 The IGIMF and its consequences (Weidner & Kroupa 2004)
IGIMF stands for Integrated Galactic Initial Mass Function
It assumes that all the stars in a galaxy are formed in embedded star clusters (Lada & Lada 2003).
Within each embedded cluster the stars are distributed according to a fixed, canonical IMF (Kroupa 2002)
The IMF within each cluster is truncated according to the total mass of the cluster (the more massive the cluster, the higher the upper stellar mass)
Stellar clusters are distributed according to a power-law -β distribution function ξ(Mcl) ∝Mcl dMcl.
Observations suggest β to be about 2.
The largest star cluster in a galaxy depends on the star formation rate (the higher the star formation rate, the higher the probability to find massive star clusters)
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The IGIMF and its consequences
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44 The IGIMF and its consequences: [α/Fe] ratios vs. mass
Recchi et al. 2009
The velocity dispersion σ is a proxy of the mass
The best fit IGIMF model (red line) reproduces the observed data better than the best fit constant IMF model (magenta line)
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Chemical feedback by the IMFs
What do we expect? In the case of lacking massive stars α-element yields should be reduced.
filled IMF truncated IMF
For the truncated IMF [O/Fe] becomes < 0; observed e.g. in dSphs.
The WSsame 2018/19 should be studied for BaCDE vs. IV Mg! 94 Steyrleithner , G.H.,et al. (2017)
45 Ba vs. Mg of MW halo stars
Various explanations of the huge Ba/Mg scatter are proposed; the formation of the halo by disrupted star clusters (no GCs) of different, but also low masses with various lack of massive stars provides a natural explanation.
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