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Star Formation in Context. Neal Evans University of Texas at Austin. Context. Individual star formation in detail Initial conditions and early evolution Formation and evolution of disks Provides the context for planet formation Massive, clustered star formation Details less accessible

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Star formation in context

Star Formation in Context

Neal Evans

University of Texas at Austin


Context
Context

  • Individual star formation in detail

    • Initial conditions and early evolution

    • Formation and evolution of disks

    • Provides the context for planet formation

  • Massive, clustered star formation

    • Details less accessible

    • Statistical results

    • Context for galaxy formation and evolution


Low mass star formation
Low-mass Star Formation

Features:

Dusty envelope

Rotation

Disk

Bipolar outflow

R. Hurt, SSC



Science goals
Science Goals

  • Complete database for nearby (< 350 pc) regions

    • Low mass star and substar formation

  • Follow evolution: starless cores to planet-forming disks

  • Coordinate with FEPS team

    • ensure complete coverage of 0 to 1 Gyr

  • Cover range of other variables

    • mass, rotation, turbulence, environment, …

    • separate these from evolution.


A typical starless core
A Typical Starless Core

L1014 distance ~ 200 pc, but somewhat uncertain.

R-band image;dust blocks stars behind and our view of what goes on inside.


Forming star seen in infrared
Forming Star Seen in Infrared

Three Color Composite:

Blue = 3.6 microns

Green = 8.0 microns

Red = 24 microns

R-band image from DSS at

Lower left.

We see many stars through the cloud not seen in R.

The central source is NOT a background star.

L1014 is forming a star

(or substar)

C. Young et al. ApJS, 154, 396


Jhk image
JHK Image

J, H, K Image of L1014

KPNO 4-m + Flamingos

J (19.7) H(20.9) K(19.4)

Huard et al. in prep.

Preliminary reduction

Faint conical nebula to north with apex on IRAC source.

BIMA peak to south likely obscures southern lobe.

Not a background source.



Lessons from l1014
Lessons from L1014

  • “Starless” cores may not be

    • Or may have substellar objects

  • Very low luminosity sources may exist

    • Must be low mass and low accretion

    • Very compact outflow detected with SMA

    • Peculiar, non-thermal radio source

  • Early, small disks are easily detected

    • Md ~ 4 x 10–4 Msun (Rd/50AU)0.5

    • Easily detected (SNR = 50–100)

    • Rd ~ 50 AU

  • Are there others?

    • About 15 candidates


Detailed studies of low mass star formation
Detailed Studies of Low Mass Star Formation

  • Isolation

  • Nearby

    • Good spatial resolution

    • Can study faint features

  • Predictive theories exist

    • n( r), v( r)

    • With caveats, L(t)

      • Allows one to calculate T(r)


We need a self consistent model
We need a self-consistent model

  • All quantities vary along line of sight

    • Dust temperature, Td( r)

      • Heating from outside, later inside

    • Gas temperature, TK( r)

      • Gas-dust collisions, CRs, PE heating

    • Density, n(r), variations predicted, observed

    • Velocity, v(r), variations predicted, observed

    • Abundance, X(r), variations predicted, observed

      • Photodissociation, freeze-out, desorption


An evolutionary model
An Evolutionary Model

  • Assume a slow approach to collapse

    • Sequence of Bonnor-Ebert spheres

      • From nc = 104 to 107 in 6 steps of factors of 3

      • Total time 1 million years

      • Time step shrinks by factor of 2 for each step in nc

      • Embedded in cloud with AV = 0.5 (or 3) mag

    • Approach a singular isothermal sphere

  • Initiate collapse at t =0

    • Inside-out collapse (Shu)

    • At each time step, calculate:

      • n (r), v(r), L, Td(r), TK(r), X(r)

      • Chemistry follows gas parcels during collapse

      • The radiation field, density, temperature change as it falls



L t from accretion contraction
L(t) from Accretion, Contraction

L(t) calculated.

First accretion.

First onto large (5 AU) surface

(first hydrostatic core).

Then onto PMS star with

R = 3 Rsun, after 20,000 to 50,000 yr. And onto disk.

Prescriptions from Adams and Shu.

Contraction luminosity and deuterium burning dominates after t ~100,000 yr.

C. Young and Evans, submitted.


Dust radiative transport
Dust Radiative Transport

Use DUSTY to compute Td( r, t). Include interstellar radiation and central heating from L(t).

Compute SED and radial

profile from DUSTY and obssphere.

C. Young and Evans, submitted


Calculate gas temperature
Calculate Gas Temperature

J. Lee et al.

2004

Use gas energetics code (Doty) with gas-dust collisions, cosmic rays, photoelectric heating, gas cooling. Calculate TK( r, t).


Calculate abundances
Calculate Abundances

Chemical code by E. Bergin

198 time steps of varying length, depending on need.

Medium sized network with 80 species, 800 reactions. Follows 512 gas parcels.

Includes freeze-out onto grains and desorption due to thermal, CR, photo effects. No reactions on grains. Assume binding energy on silicates for this case.

J. Lee et al. 2004


A closer look
A Closer Look

A few abundance profiles at t=100,000 yr.

Vertical offset for convenience

(except CO and HCN).

Big effect is CO desorption, which affects most other species. Secondary peaks related to evaporation of other species.

J. Lee et al. 2004


Calculate observables
Calculate Observables

Line profiles calculated from Monte Carlo plus virtual telescope codes. Includes collisional excitation, trapping.

Variations in density, temperature, abundance, velocity are included.

Assumes distance of 140 pc and typical telescope properties.

J. Lee et al. 2004

J. Lee et al. In prep


A closer look1
A Closer Look

Lines of HCO+ (J = 1–0 and 3–2).

Shown for four times and for different amounts of extinction in the surrounding medium.

Blue profiles are indicative of collapse.

J. Lee et al. 2004


Dynamical vs static models
Dynamical vs Static Models

Previous chemical models did not include dynamical evolution.

A fully dynamical model captures the changing density and temperature of a gas parcel.

J. Lee et al. 2005, in prep.


Abundances differ
Abundances Differ

Solid curve show result of dynamical model; dashed is static model.

In particular, the peaks at small radii from direct evaporation of the molecule are missed by static models.

J. Lee et al. 2005, in prep.


Consequences for line profiles
Consequences for Line Profiles

J. Lee et al. 2005, in prep.


Comparison to observations
Comparison to Observations

Observations of B335

Three CS transitions

Red line is from chemical model.

Evans et al. 2005, submitted


Hcn in b335 and model
HCN in B335 and Model

Evans et al. 2005, submitted


3d vs 1d dust models
3D vs 1D Dust Models

3D Dust radiative transfer:

Allows modeling of shape.

Application to L1544.

Overall results similar to 1D.

Nice constraint on internal luminosity from shape of contours.

Doty et al. 2005, MNRAS, in press


Summary so far
Summary so far

  • Beginning to develop evolutionary models

    • Self-consistent physical, chemical models

    • Get a feel for what parameters affect what

    • Starting to explore 3D models

  • Chemistry is crucial to physical modeling

    • Conclusions about dynamics from lines depend on X(r)

    • Surrounding cloud/external radiation important

  • Future work

    • Change dust opacities when ices evaporate

    • Check effect of chemistry on energetics

    • Try other dynamical models


Star formation in larger clouds
Star Formation in Larger Clouds

  • Where do stars form in large molecular clouds?

    • Early evidence indicated only in dense gas

      • Lada et al. 1991: Study of L1630

    • But surveys were incomplete

      • Need to survey at longer wavelengths

    • Large cloud surveys with c2d and COMPLETE


Perseus 12 co map
Perseus 12CO Map

The COMPLETE Team; Ridge et al.in prep.


Perseus 1mm continuum
Perseus 1mm Continuum

M. Enoch et al., in prep.


Perseus 13 co
Perseus 13CO

The COMPLETE Team; Ridge et al.in prep.


Perseus mips 24 70
Perseus MIPS (24+70)

Stapelfeldt et al. in prep.


Perseus zoom
Perseus Zoom

IRAC1 (blue), IRAC3(green, MIPS1(red)


Complementary Millimeter & Spitzer IRAC/MIPS Observations: B1 in Perseus

Bolocam 1mm

MIPS 24 m

Enoch et al. in prep.

Dark blue: 24; light blue 70 microns


Lessons from perseus
Lessons from Perseus B1 in Perseus

  • There are interesting things going on outside the famous regions

  • Large surveys needed to remove bias

  • A panchromatic view is needed

    • Molecular emission

      • Different molecules show different things

    • Dust continuum emission (across wavelengths)

    • Locations, L, etc. of forming stars

  • Much analysis remains to be done…


Disk evolution
Disk Evolution B1 in Perseus

  • Do all solar-mass stars have disks?

    • Do weak-line T Tauri stars have debris disks?

    • Are there variables besides time?

  • What are the timescales for disk evolution?

    • Formation and early evolution during collapse

    • How does the transition from accretion disks to debris disks depend on time and other factors?

  • What is the structure of disks?

  • What is the chemistry in disks?


Evidence for large inner rims in cTTs B1 in Perseus

The solid blue line (Total SED A) corresponds to the total SED when the inner rim is irradiated only by the photosphere of the central star (rim A). The solid red line (Total SED B) corresponds to the total SED when the emission from the inner rim is scaled by a factor W. W ranges from ~ 1 to ~7.

The inner rim is powered by more than the stellar photosphere

Missing source of energy?

UV radiation from the accretion shock

cgplus model (Dullemond et al. 2001)

~ 3.5

RrimA~ 0.04 AU RrimA~ 0.07 AU

Cieza et al. in prep.


Finding disks with mips
Finding disks with MIPS B1 in Perseus

Model has 0.1 Mmoon of

30 mm size dust grains

in a disk from 30–60 AU

Bars are 3 s

Model based on disks

around A stars


A new disk in cha ii
A New Disk in Cha II B1 in Perseus

  • One source not previously identified as a YSO on

    K vs. K-24 plot

  • Factor of 3 excess in 24 micron flux over stellar model indicating the presence of a disk

K.Young et al. in prep.


24m B1 in Perseusm Excesses vs Age

decay over ~ 200 Myr - many stars of all ages have no, or very little excess.


Chemistry in Disks B1 in Perseus

Robert Hurt, SSC



Studies of high mass regions
Studies of High Mass Regions B1 in Perseus

  • Many Detailed Studies

    • Ho, Zhang, Keto, …

  • Surveys

    • van der Tak et al. (2000) (14 sources)

    • Beuther et al. (2002) (69 sources)

    • Survey of water masers for CS

      • CS survey Plume et al. (1991, 1997)

        • Dense: <log n> = 5.9

      • Maps of 51 in 350 micron dust emission

        • Mueller et al. 2002

      • Maps of 63 in CS J = 5–4 emission

        • Shirley et al. 2003


Luminosity versus mass
Luminosity versus Mass B1 in Perseus

Log Luminosity vs. Log M

red line: masses of dense

cores from dust

Log L = 1.9 + log M

blue line: masses of GMCs

from CO

Log L = 0.6 + log M

L/M much higher for dense cores than for whole GMCs.

Mueller et al. (2002)


Linewidth versus size
Linewidth versus Size B1 in Perseus

Correlation is weak.

Linewidths are 4-5 times larger than in samples of lower mass cores.

Massive clusters form in regions of high turbulence, pressure.

Shirley et al. 2003


Cumulative mass function
Cumulative Mass Function B1 in Perseus

Incomplete below 103 Msun.

Fit to higher mass bins gives slope of about –0.93.

Steeper than that of CO clouds or clumps (–0.5 on this plot).

Similar to that of clusters, associations (Massey et al. 1995) in our Galaxy and in Antennae (Fall et al. 2004).

Shirley et al. 2003


Hints of dynamics
Hints of Dynamics B1 in Perseus

A significant fraction of the massive core sample show self-reversed, blue-skewed line profiles in lines of HCN 3-2.

Of 18 double-peaked profiles, 11 are blue,

3 are red.

Suggests inflow motions of overall core.

Vin ~ 1 to 4 km/s over radii of 0.3 to 1.5 pc.

J. Wu et al. (2003)


Low mass vs high mass
Low Mass vs. High Mass B1 in Perseus

  • Low Mass star formation

    • “Isolated” (time to form < time to interact)

    • Low turbulence (less than thermal support)

    • Slow infall

    • Nearby (~ 100 pc)

  • High Mass star formation

    • “Clustered”

    • Time to form may exceed time to interact

    • Turbulence >> thermal

    • Fast infall?

    • More distant (>400 pc)


High vs low early conditions
High vs. Low Early Conditions B1 in Perseus

n( r) = nf (r/rf)–p ; rf = 1000 AU


Massive cores gross properties
Massive Cores: Gross Properties B1 in Perseus

  • Massive, Dense, Turbulent

    • Mass distribution closer to clusters, stars than GMC

    • Much more turbulent than low mass cores

    • Similar overall power law shape

    • About 100 times denser

    • Linewidths about 16 times wider


Star formation in galaxies
Star Formation in Galaxies B1 in Perseus

“Schmidt Law”

Kennicutt Relations

SSFR = A SN(gas);

N = 1.4 to 2.4

Log SSFR

Log S(gas)

Kennicutt, 1998, ARAA, 36, 189


Dense gas in galaxies
Dense gas in galaxies B1 in Perseus

  • CO detected in many galaxies

  • Increasingly, HCN can be seen in galaxies

  • Star formation rate (LFIR) correlates better with HCN than with CO


L B1 in PerseusIR

HCN

CO

Gao & Solomon 2004


Connecting to galaxy evolution
Connecting to Galaxy Evolution B1 in Perseus

  • Increasing evidence for rapid formation of some galaxies

    • Detection of many submillimeter galaxies

      • Intense starbursts at z ~2-3

      • Intense CO, dust, solar metals in QSO at z = 6.4

  • Are massive dense cores a model for starbursts?

    • L/M much higher than for GMCs as a whole

    • L/Mdust ~ 1.4 x 104 Lsun/Msun ~ high-z starbursts

    • L/L(HCN) similar to starbursts

  • Starburst: all gas like dense cores?


Coming attractions
Coming Attractions B1 in Perseus

SOFIA 2005

SMA, CARMA, eVLA, LMT, GBT, APEX, ASTE, JCMT, CSO, …

Spitzer 2003

Herschel 2007

SAFIR ~2015

JWST 2011

LSAT 2012


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