���1

Kilo-parsec scale simulations of star formation in disc galaxies

Sven Van Loo !

Harvard-Smithsonian Center for Astrophysics

!Collaborators:

JonathanTan (Florida) Mike Butler (Zurich)

Sam Falle (Leeds) !!

Sapporo 17.02.2014

���2

!

• Observations of molecular clouds in galaxies

• Small-scale models of galactic discs

• Magnetic fields

• Conclusion

Outline:

���3

Observations of star formation in disc galaxies

In galaxies stars do not form everywhere, but are associated with Giant Molecular Clouds

Whirlpool galaxy/M51 (NASA/ESA, Beckwith & Hubble Heritage Team) Grand design spiral galaxy

���3

Observations of star formation in disc galaxies

In galaxies stars do not form everywhere, but are associated with Giant Molecular Clouds

Whirlpool galaxy/M51 (NASA/ESA, Beckwith & Hubble Heritage Team) Grand design spiral galaxy

CO emission (Koda et al. 2009)

���3

Observations of star formation in disc galaxies

In galaxies stars do not form everywhere, but are associated with Giant Molecular Clouds

Whirlpool galaxy/M51 (NASA/ESA, Beckwith & Hubble Heritage Team) Grand design spiral galaxy

CO emission (Koda et al. 2009)

And then only in the densest regions of the clouds.

W3 (ESA/PACS&SPIRE consortia, Rivera-Ingraham & Martin)

���4

Observational information Combine the information from own Galaxy with others

(Bigiel et al. 2008)

→ Star formation efficiency is low Cloud size/velocity line width relation → Suggestion of energy cascade

(Heyer & Brunt 2004)

���4

Observational information Combine the information from own Galaxy with others

(Bigiel et al. 2008)

→ Star formation efficiency is low Cloud size/velocity line width relation → Suggestion of energy cascade

(Heyer & Brunt 2004)

Same origin? (Krumholz, McKee & Tumlinson 2009)

���5

Origin of turbulence

3 ways of understanding GMC formation and evolution !! • Bottom-up processes: agglomeration due to cloud-cloud collisions,

compression induced in converging flows, compression of extended, but finite structures by shocks (e.g. supernova explosions, supershells) • Top-down processes: large-scale gravitational, thermal and magnetic instabilities in the differentially rotating disc of a galaxy • Continuous evolution: balance of cloud mergers, destruction, dispersion, ... maintaining a reservoir of molecular gas

!!Current observations favour the top-down model for GMC formation and continuous evolution (Koda et al. 2009), but the relative importance of top-down and bottom-up mechanisms may depend on the galactic environment (Dobbs 2008) !!!!!!

���6

Large to small scale simulations

In global simulation only atomic cooling down to 300K (a=1.8 km/s), self-gravity and external gravitational potential. Follow kpc extract from global simulation to sub-pc scales and add additional physics (e.g. molecular cooling, star formation prescription, magnetic fields).

6 clouds at 4.25 kpc with 0.8-7.5x106 M☉ and velocity dispersions of ~10 km/s.

(Tasker & Tan 2009)

���7

Small scale simulations of galaxy simulations

~ 8 pc ~ 0.5 pcSame physics, but higher resolution (follow up to 10 Myr)

• Mass fraction of gas in clouds remains similar, but gas is redistributed to higher densities • Cloud properties (mass, velocity dispersion) in both resolutions similar • More filamentary clouds

���8

Molecular coolingInclude treatment for molecular cooling (down to 5K) without radiative transfer !From the numerical simulation with Sanchez-Salcedo cooling → density/extinction relation Simplifies parameter space (n, T, Av) for calculating heating and cooling rates with Cloudy

Cooling rates can be used to calculate emissivities and line profiles in the simulations (see Ben Wu’s talk)

���8

Molecular coolingInclude treatment for molecular cooling (down to 5K) without radiative transfer !From the numerical simulation with Sanchez-Salcedo cooling → density/extinction relation Simplifies parameter space (n, T, Av) for calculating heating and cooling rates with Cloudy

Cooling rates can be used to calculate emissivities and line profiles in the simulations (see Ben Wu’s talk)

���9

Molecular cooling

• Similar results as previous cooling functions → thermal support of the clouds is not important • Clouds maintain a constant density distribution for several free-fall times → support provided by e.g. galactic shear cascades down into clouds

!• Column densities are too high compared (x2-3) to observations → need to remove gas and form stars

(Van Loo, Butler & Tan 2013)

���9

Molecular cooling

• Similar results as previous cooling functions → thermal support of the clouds is not important • Clouds maintain a constant density distribution for several free-fall times → support provided by e.g. galactic shear cascades down into clouds

!• Column densities are too high compared (x2-3) to observations → need to remove gas and form stars

Rosen & Bregman cooling

(Van Loo, Butler & Tan 2013)

���9

Molecular cooling

• Similar results as previous cooling functions → thermal support of the clouds is not important • Clouds maintain a constant density distribution for several free-fall times → support provided by e.g. galactic shear cascades down into clouds

!• Column densities are too high compared (x2-3) to observations → need to remove gas and form stars

Rosen & Bregman cooling

0 2 4 6 8 10Time [Myr]

0

0.2

0.4

0.6

0.8

1

Mas

s fra

ctio

n

1x10-4

2x10-4

3x10-4

4x10-4

5x10-4

Volu

me

fract

ion

(Van Loo, Butler & Tan 2013)

Cloud gas: > 102 cm3

Clump gas: > 105 cm3

Volume gas: > 102 cm3

���10

Star formationInclude stochastic star formation prescription: Star particles form

•when n > 105 cm-3 and T < 3000K •with mass M=εff ρg Δx3 Δt/tff with εff =0.02

but only with Mmin=100M⊙ to avoid too many particles of low mass. If mass is below Mmin then particle is formed with probability of M/Mmin !!

(Van Loo, Butler & Tan 2013)

���10

Star formationInclude stochastic star formation prescription: Star particles form

•when n > 105 cm-3 and T < 3000K •with mass M=εff ρg Δx3 Δt/tff with εff =0.02

but only with Mmin=100M⊙ to avoid too many particles of low mass. If mass is below Mmin then particle is formed with probability of M/Mmin !!

• Lower column densities • Star formation rate too high

(Van Loo, Butler & Tan 2013)

���11

Analysis of filaments

Study density structure and kinematics of filaments in even higher resolution simulations ~ 0.13 pc (Butler et al., in prep.).

Compare with observations: !!• Probability Density Function of

filaments are log-normal with high-end power-law tail (Kainulainen et al. 2009) !

• Equally spaced dense clumps (e.g Nessie Nebula; Jackson et al. 2010) !

• Filament shape: Plummer-profile in column density and density? (Arzoumanian et al. 2012).

See Mike Butler’s poster

Magnetic fields are important from galaxy down to core scale

���12

Magnetic field observations

M51 (Fletcher et al 2010)

NGC1333 IRAS4 (Girart et al. 2006)

Magnetic fields are important from galaxy down to core scale

���12

Magnetic field observations

M51 (Fletcher et al 2010)

NGC1333 IRAS4 (Girart et al. 2006)

M33 (Li & Henning 2011)

Magnetic fields anchor the clouds

���13

Magnetic fields in local simulations

0μG 10μG

Include a uniform magnetic field along the shearing direction (y-direction) and vary the magnitude (Galactic field ~ 6μG)

40μGresolution ~ 2 pc

���13

Magnetic fields in local simulations

0μG 10μG

Include a uniform magnetic field along the shearing direction (y-direction) and vary the magnitude (Galactic field ~ 6μG)

40μGresolution ~ 2 pc

Internal structure of the cloud changes, i.e. less fragmentation.

���13

Magnetic fields in local simulations

0μG 10μG

Include a uniform magnetic field along the shearing direction (y-direction) and vary the magnitude (Galactic field ~ 6μG)

40μGresolution ~ 2 pc

���14

Internal structureFocus on single cloud

• Less fragmentation with stronger magnetic field (Hennebelle et al. 2011) • Magnetic field changes direction within the cloud, although the large-scale field remains orientated along shearing direction. → clouds anchored similar as in M33 (Li & Henning 2011) !• Filaments axis either perpendicular or parallel to magnetic field similar to observations (Li et al. 2013). !

���14

Internal structureFocus on single cloud

• Less fragmentation with stronger magnetic field (Hennebelle et al. 2011) • Magnetic field changes direction within the cloud, although the large-scale field remains orientated along shearing direction. → clouds anchored similar as in M33 (Li & Henning 2011) !• Filaments axis either perpendicular or parallel to magnetic field similar to observations (Li et al. 2013). !

���14

Internal structureFocus on single cloud

• Less fragmentation with stronger magnetic field (Hennebelle et al. 2011) • Magnetic field changes direction within the cloud, although the large-scale field remains orientated along shearing direction. → clouds anchored similar as in M33 (Li & Henning 2011) !• Filaments axis either perpendicular or parallel to magnetic field similar to observations (Li et al. 2013). !

(Li et al 2013)

���15

Global evolution

Mass properties of the clouds

•Difference in cloud mass due to magnetically dominated diffuse gas not being accreted onto clouds. !

• Clump mass unchanged: internal structure determined by turbulent gas motions. Clouds initially supercritical. !

Cloud gas: > 102 cm3

Clump gas: > 104 cm3

���15

Global evolution

Mass properties of the clouds

•Difference in cloud mass due to magnetically dominated diffuse gas not being accreted onto clouds. !

• Clump mass unchanged: internal structure determined by turbulent gas motions. Clouds initially supercritical. !

Cloud gas: > 102 cm3

Clump gas: > 104 cm3

Stars form in the mass reservoir of clumps, so no reduction in SFR.

���16

Critical transitionMagnetic field-density distribution

• Maximum magnetic field values: B ∝ n2/3

→ isotropic contraction in ideal MHD !• Simulation distribution shows a different slope at the high end → also flow along the magnetic field lines !

(Crutcher et al 2010)

Supercr

iticalSub

critical

���17

Conclusions

• Clouds do not collapse as they maintain a constant density distribution. !• Self-gravity leads to too much star formation !

• Magnetic fields do not seem to suppress the mass accumulation in clumps !

• Magnetic fields do suppress fragmentation of the clouds !