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University of Heidelberg, Center for Astronomy

Giant Star-Forming Regions

Dimitrios A. Gouliermis

Lecture #12 Star Formation

Giant Star-Forming Regions Gouliermis & Klessen

Schedule of the Course

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Lect. 1 19-Oct-2012 Course Overview Motivation for the Course/Schedule; Overview of Physical Processes in HII Regions; Classification of HII regions

Lect. 2 26-Oct-2012 Introduction to the Physics of the ISM I Phases of the ISM; Transitions; Introduction to cooling mechanisms

Lect. 3 2-Nov-2012 Introduction to the Physics of the ISM II Atomic Transitions; Gas Cooling; Collisional Excitation

Lect. 4 9-Nov-2012 Introduction to the Physics of the ISM III Gas Heating; Photo-ionization; Photo-electric heating; PAHs

Lect. 5 16-Nov-2012 Interstellar Dust Composition; Spectral Features; Grain Size Distributions; Extinction

Lect. 6 23-Nov-2012 Physical Processes in HII Regions I Radiative Processes; Photo-ionization & Recombination of hydrogen; Photoionization Equilibrium

Lect. 7 30-Nov-2012 Physical Processes in HII Regions II Heating and Cooling of HII Regions; Strömgren Theory; Forbidden lines and Line Diagnostics

Lect. 8 7-Dec-2012 Physical Processes in HII Regions III T & n measurements in HII Regions; Secondary Ionization; Dielectronic Recombination; Free-free emission

Lect. 9 14-Dec-2012 Photodissociation regions (PDR) Energy Balance; Dissociation of Molecular Hydrogen; Structure; Observations & Spectral features

Lect. 10 11-Jan-2012 Stellar Feedback I Introduction; The Physics of Shocks; Super-Nova Evolution; SNR

Lect. 11 18-Jan-2012 Stellar Feedback II Expansion of HII regions; Wind-Blown Bubbles; Super-bubbles/Triggered Star Formation

Lect. 12 25-Jan-2012 Star Formation Triggered SF; Protostars and YSOs; PMS evolution


Star Formation in Giant HII Regions (Observer’s view)

In this Lecture •  Star Formation & Triggered SF •  Young Stellar Objects •  Pre–Main-sequence Evolution •  T Tauri and Herbig Ae/Be Stars

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Star Formation in Giant HII Regions (Observer’s view)

Suggested Bibliography –  S. W. Stahler & F. Palla, The Formation of Stars, WILEY-VCH

GmbH (2004) –  F. H. Shu, F. C. Adams & S. Lizaro, Star Formation in molecular

clouds - Observation and theory, Annu. Rev. Astron. Astrophys. 25, pp. 23–81 (1987)

–  Ch. J. Lada, The Formation of Low Mass Stars: An Observational Overview, in “The Origin of Stars and Planetary Systems”, p. 143, Kluwer (1999)

–  Bo Reipurth (Ed.), Handbook of Star Forming Regions, Vol. I & II, ASP Monograph Publications (2008)

–  F. Palla, The Evolution of Pre-Main-Sequence Stars, in “The Origin of Stars and Planetary Systems”, p. 375, Kluwer (1999)

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Star Formation

•  The formation of stars is governed by the complex interplay between gravitational attraction in molecular clouds and supersonic turbulence, thermal pressure, and magnetic fields, “stored” in the interstellar matter.

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Christopher F. McKee and Eve C. Ostriker “Theory of Star Formation” Annu. Rev. Astron. Astrophys. 45 pp. 565–687 (2007)

Ralph E. Pudritz “Clustered Star Formation and the Origin of Stellar Masses” Science 295, Issue 5552, pp. 68-76 (2002)


Star Formation (Cores formation) •  SF initiates when massive (~107 M⊙) bound structures condense out

of the diffuse ISM as a result of gravitational instabilities. •  Internal turbulence inherited from the diffuse ISM in combination

with self-gravity causes fragmentation into GMCs and clumps. •  Turbulence within GMCs is highly supersonic and it imposes a log-

normal distribution of densities. This structure is hierarchical. •  Turbulence damps in about one crossing time, and it is not under-

stood how and for how long it can be maintained in the universal level observed in GMCs.

•  Spatially defined structures within GMCs tend to have internal velocity dispersions that increase with size as σ ∝ l0.5.

•  The densest regions become self-gravitating cores with masses that are typically of the order of the Bonnor-Ebert mass.

•  These cores are frequently clustered, owing to the dominance of large scales in the turbulent flow. Their formation rate depends on the turbulent properties of the GMC.

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Star Formation (Core-collapse) •  Magnetically supercritical cores undergo collapse, first

becoming strongly stratified internally. •  Continued accretion occurs if the surrounding ambient medium

has a sufficiently low level of turbulence. •  Collapse leads to the formation of a rotating disk interior to an

accretion shock; significant magnetic flux is lost in this process. •  Disks accrete owing to processes that transport angular

momentum outward through gravitational and magnetic stresses. •  Powerful winds are magnetocentrifugally driven from the circum-

stellar disks. The inner wind becomes collimated into a jet-like flow. •  A stratified disk wind sweeps up much of the ambient gas into a

massive molecular outflow. This reduces the net SFE to ~1/3. •  Massive stars form from cores that are considerably more massive

than a Bonnor-Ebert mass, and are most likely highly turbulent. •  Radiation pressure can be overcome by the disk formation, proto-

stellar outflows, and hydrodynamic instabilities in the accreting gas. WS 2012 - 2013 Lecture 12 7


Star Formation (GMC destruction)

•  Massive, luminous stars ionize their surroundings into HII regions. The expansion of these regions at ~10 km s−1 contributes to the large-scale turbulent power.

•  This process can unbind GMCs within a few crossing times. By the time of their destruction, GMCs may have lost much of their original mass by photoevaporation.

•  The destruction of GMCs returns almost all of their gas to the diffuse phase of the ISM, with a mean SFE over the cloud lifetime of ~5%.

•  This low efficiency is a consequence of the small fraction of mass compressed into clumps dense enough to resist turbulence.

•  The return of GMC gas to the diffuse ISM completes the cycle of star formation, which then begins anew.

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The Recycling Process of the ISM

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Credit: Roland Diehl in "Astronomy with Radioactivities VII” workshop


Introduction to Young Stellar Objects •  YSOs are stars under formation. They are active, with

dynamic circumstellar disks, bipolar outflows, and infalling dusty envelopes.

•  YSOs are studied in detail in nearby GMCs •  Observations are affected by extinction, which

decreases with wavelength. Embedded sources - visible only from the NIR to mm and considered to be very young Revealed sources - visible at optical-NIR wavelengths, and considered more evolved (~1Myr)

•  YSOs are associated with early stellar evolution phenomena like masers, Herbig-Haro objects, and proplyds.

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Local SF Regions •  Typical well-studied

regions: –  Taurus-Auriga –  Perseus –  Orion

•  For a complete account see “Handbook of Star Forming Regions” edited by Bo Reipurth

•  Most of info about YSOs comes from these regions. WS 2012 - 2013 Lecture 12 11


Characteristics of YSOs •  YSOs consist of several components :

–  Cloud cores (cocoons) –  Accretion Disks –  Ouflows: jets and swept up molecular gas

•  YSO components are observed from optical to mm bands. •  They vary in size from a few solar radii to 10,000 AU. •  Spatially resolved observations at nearby regions (Taurus)

require arcmin down to milli-arcsec spatial resolution. •  Signposts of ongoing star formation:

–  Embedded IR-bright Sources –  Dark clouds and dense cores –  Young Stellar Clusters and Associations

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The ONC in the near-IR

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A color composite mosaic image of the Orion Nebula Cluster, based on images from the near-IR camera ISAAC on the ESO Very Large Telescope.

The famous Trapezium stars are seen near the center, and the photo also shows the associated cluster of about 1000 stars, about a million years old. Credit: ESO/Mark McCaughrean.


Late Accretion Phase – Bipolar Outflow (Pre—Main-Sequence star)

Optically Thin Disk

YSO Evolution Sequence (Low-Mass Star Formation Scenario)

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Main Accretion Phase (Embedded Protostar)

Optically Thick Disk Literature: Shu, Adams & Lizaro 1987, ARA&A, 25, p.23 Lada 1987, in “Star-Forming Regions”, IAU Symp. 115, p.1


Spectral Energy Distributions of YSOs •  Pioneering SED classifications, based on the fact that YSOs

are obscured by circumstellar material that reradiates in IR –  Lada and Wilking (ApJ 287 610 1984) –  Adams, Lada and Shu (ApJ 312 781 1987).

•  The reprocessing agent is a circumstellar disk •  YSOs are classified into the following classes:

–  Class 0 - earliest phase of Class 1 - just after collapse –  Class 1 - deeply embedded sources, with positive IR spectral index,

emitting in the mm and sub-mm –  Class 2 - revealed sources with near-IR and mid-IR excess –  Class 3 - small near-IR excess

•  This classification is based on the continuum (dust) spectrum.

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SEDs of YSOs Classes

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Class 0 Class I

Class II Class III

Cold Black Body

Disk

StellarPhotosphere

Main accretion phase (10 yr)Late accretion phase (10 yr)

- bipolar out!ow

Optically thick disk (10 yr) Optically thin disk (10 yr)

4

5

6 7

Literature: Lada 1999, in “The Origin of Stars and Planetary Systems”, (NATO ASI) Shown Models by: Robitaille et al. 2006, Astroph. J. Suppl. 167, p.256


SED Classification of YSOs •  Charlie Lada in 1987 introduced the classification criteria

based on the values of SED’s spectral index:

•  He proposed three YSO classes •  The spectral index is calculated in the interval λ = 2.2 – 20

µm (near- and mid-IR). •  In 1993 Phillippe André et al. (Astroph. J. 406, p.122)

discovered Class 0; objects with strong submillimeter emission, but very faint at λ < 20 µm.

•  Later Thomas Greene et al. in 1994 (Astroph. J. 434, p.614) added the class of "flat spectrum" class sources.

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€

α =d log λFλ( )d log λ( )


SED Classification of YSOs •  Class 0 sources – undetectable at λ < 20 µm. •  Class I sources have α > 0.3 •  Flat spectrum sources have –0.3 < α < 0.3 •  Class II sources have –1.6 < α < –0.3 •  Class III sources have α < –1.6

Online SED Fitter: http://www.hyperion-rt.org/ WS 2012 - 2013 Lecture 12 18

This classification scheme roughly reflects an evolutionary sequence. The most deeply embedded Class 0 sources are considered to evolve towards Class I stage dissipating their circumstellar envelopes. Eventually they become optically visible on the stellar birthline as pre–main-sequence stars.


Pre–Main-Sequence (PMS) Evolution •  The PMS phase in the evolution of stars with

masses up to ~6 M⊙ corresponds to the time between the gravitational core collapse, which forms the protostar (on the birthline), and the ignition of hydrogen in the formed star, placing it on the Zero-Age Main Sequence (ZAMS).

•  During this evolutionary phase, the observed radiation from the star is affected significantly by surface activity, and by circumstellar disks of dust and gas, formed by matter infalling during the collapse of the rotating core.

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Artist's impression of a T Tauri star

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PMS vs MS Stellar Structure

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From Stahler & Palla (2004)


PMS Stellar Evolution •  As the PMS contraction begins, the star is no longer buried within an

opaque dust cloud. Radiation can emanate freely from its surface layers. •  The contraction of PMS stars decelerates with time. •  Low-mass stars grow dimmer as they contract. •  Kelvin-Helmholtz time: The characteristic time scale for PMS

contraction:

•  Star’s effective temperature during PMS evolution is nearly constant, especially for subsolar masses. Therefore, since

•  It follows that tKH itself varies as L*−3/2 for a star of fixed mass.

•  Birthline is the locus in the HRD of PMS stars with protostellar radii. •  The ignition of hydrogen ends contraction for normal stars.

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€

tKH =GM∗

2

R∗L∗, with dR∗

dt= −n1

R∗tKH

, n1 ≈1

€

L∗ = 4π R∗2σB Teff

4 ⇒ dL∗dt

= −n2L∗tKH

, n2 ≡ 2n1


PMS Stellar Evolution •  0.1 M ≤ M ≤ 1 M: for low-mass protostars, deuterium burning occurs

near the center for temperatures exceeding 106 K. Deuterium acts as an effective thermostat, preventing central temperature from rising to ≥106K appropriate for hydrogen burning. Having consumed most of its deuterium during accretion, the star begins the contraction phase.

•  1 M ≤ M ≤ 2.5 M: Fully Convective Stars due to surface cooling. The evolution is undisturbed by the fusion of residual deuterium.

•  2.5 M ≤ M ≤ 4 M: Partially Convective Stars that undergo thermal relaxation and nonhomologous contraction, from which stems almost entirely their luminosity.

•  4 M ≤ M ≤ 8 M: Fully Radiative Stars. The strength of the gravitational pull prevents a further increase of the radius. They contract homologously under their own gravity.

•  Central hydrogen burning is delayed until ~ 6 M. The ZAMS is immediately reached for stars with mass larger than ~ 8 M.

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Theoretical vs Observational PMS HRD

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HR diagram of the Taurus-Auriga SF Region (Stahler 1988, ApJ, 332, 804) and the Orion Nebula Cluster (From Stahler & Palla 2004). Birthline, ZAMS, and selected pre–main-sequence tracks are shown. Stellar masses for the tracks are given in solar units.


T Tauri Stars (TTSs)

•  Original Optical Identification was made by A.H. Joy in 1945 (Astroph. J. 102, p. 168).

•  Defined after the prototype T Tau. •  Spectral Types F5 to M5. •  Emission lines patterns that resemble the Solar

chromosphere, but they are much stronger. •  Highly Irregular variability. •  Ambartsumian proposed in 1957 that TTSs are

YSOs.

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Characteristics of T Tauri Stars

•  Optical emission lines from chromospheric heating.

•  Emission in X-rays, UV due to hot coronae and magnetic activity.

•  Periodic fluctuations indicating rotating star (cool & hot) spots.

•  Optical variability because of accretion. •  Excess broadband flux in U, Hα and IR due to

circumstellar disks. •  Molecular outflows, winds and accretion.

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Identification Criteria of T Tauri Stars

•  Optical Spectroscopy •  Balmer H, Ca II, Fe I, Fe II, Ti II, He I emission lines •  Li I λ6707 absorption

•  UV Spectroscopy •  Mg II, Fe II, He II, C I-IV, S II-IV, N II-V emission lines

•  Photometric U-, Hα- and IR-excess •  Timely dependent brightness (variability) •  X-ray visibility

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Literature: Appenzeller & Mundt 1989, Astron. & Astroph. Rev. 1, p.291 Bertout 1989, Ann. Rev. Astron. & Astroph. 7, 351 Feigelson & Montmerle 1999, Ann. Rev. Astron. & Astroph. 37, 363


Spectroscopy of T Tauri Stars

•  Hα is a primary line diagnostic, associated with accretion •  Classical TTS – EW(Hα) > 50Å •  Weak-lined TTS – EW(Hα) < 50Å

•  The strongest lines are usually OI, Hα and SII WS 2012 - 2013 Lecture 12 28

Kenyon et al. 1998, Astron. J. 115, p. 2491

WTTS CTTS


Accretion in TTS •  Hα line in extreme Class I TTS is very broad with

a shape indicative of strong flows (accretion and winds).

•  By definition WTTS have little or no line-emission and do not accrete.

•  This does not mean that it was always so. They probably still have disks around them.

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The Optical Orion Nebula Cluster HRD

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Hillenbrand 1997, Astron. J. 113, p. 1733

•  Pre–main-sequence stars are dated based on their HRD positions in terms of models.

•  However, there are characteristics that bias these positions. See reviews –  Gouliermis 2012, Space Science Reviews 169, p.1 –  Preibisch 2012, Research in Astron. Astrophys. 12, p.1


The updated Optical HRD of ONC

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Better-constrained age determination of the PMS populations. From Da Rio et al. 2010, Astroph. J. 722, p. 1092


Mass Accretion •  An important effect to consider is the occurrence of residual mass

accretion from a circumstellar disk during the PMS phase. •  Accretion alters the total (star+disk) luminosity, the evolutionary

path, and hence the age estimate of a star. •  Excess in near-IR emission from a PMS star is interpreted as

being due to the presence of an active circumstellar disk (Hartmann et al. 1997, ApJ 475, 770).

•  The accreted material carries liberated gravitational energy at the equivalent luminosity

WS 2012 - 2013 Lecture 12 32

€

Lacc =GM∗

˙ M acc

R∗

where is the mass accretion rate •  Some fraction of this power may be liberated in UV generated on

impact of an accretion shock at the stellar surface .

€

˙ M acc = dM∗ /dt


UV Spectra of T Tauri Stars

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•  CTTS have strong UV excess emission. •  They display a Balmer continuum jump at 3647 Å. •  UV is believed to arise from an “accretion shock” onto the stellar surface.

Extracted UV-excess spectra of TTSs from Gullbring et al. ApJ 492 323 1998


Mass Accretion Rate •  Excess emission decreases the depth of the stellar photospheric

absorption lines, compared to a nonaccreting spectrum. •  The accretion rate can be determined from this “veiling” of

stellar UV absorption lines. •  Measurements of lines at different wavelengths yield the ratio of hot

excess emission to stellar photospheric emission as a function of λ. •  Gullbring et al. (1998, ApJ 492 323) established observationally an

empirical relation between the emission excess in the Johnson U-band filter and the total accretion luminosity, Lacc:

•  Gullbring et al. find values 10–9 - 10–7 M/yr in Taurus-Auriga. •  Robberto et al. (2004, ApJ, 606, 952) find 10–12 – 10–8 M/yr in

Orion.

WS 2012 - 2013 Lecture 12 34 €

log LaccL⊗

⎛

⎝ ⎜

⎞

⎠ ⎟ =1.09−0.018

+0.04 logLU ,excL⊗

⎛

⎝ ⎜

⎞

⎠ ⎟ + 0.98−0.07

+0.02


Near-IR Imaging of PMS stars

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Near-IR excess visualized in the Color-color and color-magnitude diagrams.

Hillenbrand et al. 1995, Astron. J. 109, p. 280


Herbig Ae/Be (HAeBe) Stars

•  They are more massive analogues of TTSs •  Originally identified by G. H. Herbig in 1960. •  Stars of A and B spectral types. •  Their IR-excess starts already at 1-2 µm. •  Strong mid- and far-IR emission indicates the

presence of (quasi-spherical) dusty shells in contrast to TTSs.

•  Debris Disks studies suggest that circumstellar mass declines with age as M ∝ t –2.

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SEDs of HAeBe Stars •  Comprehensive studies of the

SEDs of HAeBe stars are performed by Lynne Hillenbrand and collaborators (e.g., Hillendrand et al. 1992, ApJ, 397, 613)

•  They classify HAeBe stars in two groups:

–  Group I: Declining SED with λ –  Group II: Rising SED with λ

In the figure squares indicate observed fluxes, circles extinction corrected data. Note the onset of an IR excess already at λ~1–2 µm.

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AB Aur [Group I]

PV Cep [Group II]

Suggested Review: Waters & Waelkens, Annu. Rev. Astron. Astrophys. 36, 233 (1998)


HAeBe Stars Evolution

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Suggested evolutionary scenario: (a) embedded source; (b) single dust disk causing IR excess; (c) double dust disk; (d) single dust disk causing far-IR excess; (e) evolved dust disk. From Malfait et al. Astron. & Astroph. 331, 211 (1998).


Spectroscopy of HAeBe Stars

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Being 10-100 times more luminous than TTSs, Herbig stars are observed in mid- and far-IR wavelengths from space, providing useful information on the solids and large molecules near and around YSOs.

In the figure are shown PAH emission features at 6.2 and 7.7 µm, and 11.3 and 12.7 µm (with error bars), after continuum subtraction, obtained with Spitzer/IRS. From Sloan et al. ApJ 632 956 (2005).


Extragalactic Young PMS Clusters

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Credit: NASA/ESA/D. Gouliermis

Gouliermis et al. 2007 ApJL, 665, 27 Da Rio, Gouliermis & Henning 2009 ApJ, 696, 528 Da Rio, Gouliermis & Gennaro 2010 ApJ, 723, 166


Triggered SF Mechanisms

Category 1 •  Pre-existence of molecular condensations or

inhomogeneities. –  Radiation-driven Implosion of pre-existing molecular

globules. •  Bright Rims (Lefloch et al. 1994; Miao et al. 2006) •  Pillars and Evaporating Gaseous Globules (EGGs; Hester et

al. 1998)

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Processes that trigger SF at the borders of HII regions


Triggered SF Mechanisms

Category 2 •  Formation of molecular condensations out of a

nearly homogeneous medium. –  Instabilities of the ionization front

•  Hydrodynamical Instabilities (García-Segura & Franco 1996; Mizuta et al. 2006)

–  Collect & Collapse process •  Due to the supersonic expansion of an HII region

(Elmegreen & Lada 1977; Whitworth et al. 1994; Hosokawa & Inutsuka 2005)

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Processes that trigger SF at the borders of HII regions


Radiatively-driven Implosion: EGGs

WS 2012 - 2013 Lecture 12 43 Hester & Desch, ASP Conference Series 341, (2005)


Evaporating Gaseous Globules

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1)  Radiation from a massive star drives an ionization front into surrounding molecular gas.

2)  The ionization front (plus winds and previous SNe) drive a shock, triggering collapse of molecular cores.

3)  ~100,000 years after triggered collapse, the ionization front overruns the core, forming an EGG.

4)  EGGs evaporate in ~10,000 years, exposing the disk. The evaporating disk is a proplyd.

5)  In ~10,000 years, disks erode to ~50 AU. Disk evaporation ends, leaving a protostar and bare protoplanetary disk.

6)  The massive star goes supernova, injecting newly synthesized elements into surrounding disks. From Hester & Desch (2005)


Collect & Collapse around Bubbles

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Zavagno et al. 2006, Astron. & Astroph., 446, 171

Schematic view of a spherical expanding HII region and of its neutral environment. Processes (1) and (2) represent small- and large-scale gravitational instabilities due to the collect & collapse mechanism. From Deharveng et al. 2010, A&A, 523, A6.


Collect & Collapse •  An HII region expands hypersonically into the

ambient ISM •  Neutral material accumulates between the ioni-

zation front and the shock front that precedes it. •  Shell may be very massive and later it collapses. •  Several massive star/cluster-forming fragments

are formed at the periphery of the HII region.

•  Proposed by Elmegreen & Lada, 1977, ApJ 214, 725. •  Analytic formulation: Whitworth et al. 1994, MNRAS 268, 291. •  Examples: Sh2-104 (Deharveng et al. 2003)

RCW 79 (Zavagno et al. 2006) WS 2012 - 2013 Lecture 12 46


Collect & Collapse

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Dale, Bonnell & Whitworth 2007, Mon. Not. Royal Astron. Soc. 375, p. 1291


Case Study: Sh2-104

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Deharveng et al. A&AL 408, 25 (2003)

12CO(2–1)

MSX (6.8–10.8 µm ) + DSS2 (Hα) + 1.46 GHz

Hosokawa & Inutsuka ApJ 623, 917 (2005)


Topics Not Covered

•  Massive Star Formation and Evolution –  Nucleosynthesis in Massive Stars –  Spectral Features –  Rotation, Binarity

•  Stellar Clusters & Associations –  Stellar Content –  Dynamical Status & Evolution

•  The Stellar Initial Mass Function –  Universallity –  Connection to Molecular Core MF

WS 2012 - 2013 Lecture 12 49

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