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Sci TR-23

Updates: 04.05.2012 (Nazım)

2.3 ISM and Star Formation

Planetary Nebulae (PNe)

A planetary nebula (PN) is an emission nebula consisting of an expanding glowing shell of ionized gas ejected during the asymptotic giant branch phase of certain types of stars late in their life. Planetary nebulae are amongst the most photogenic and complex of celestial phenomena but are also amongst the most important to properly understand. They are important probes of nucleosynthesis processes, mass-loss physics and Galactic abundance gradients, and, because their progenitor stars dominate all stars above one solar mass, PN are responsible for a large fraction of the chemical enrichment of the interstellar medium, including the seeding of pre-biotic carbon between the stars.

Furthermore, their rich emission-line spectra enable detection to large distances. These emission lines allow the determination and analysis of chemical abundances and permit the estimation of shell expansion velocities and ages, and so probe the physics and timescales of stellar mass loss (e.g. Iben 1995). The measured radial velocities can trace the kinematic properties of observed PN, enabling us to decide if they belong to a relatively young or old stellar population. The kinematic properties of PN in galaxy halos also give strong constraints both on the mass distributions and formation processes of giant elliptical galaxies (e.g. Bekki & Peng 2006; |Douglas et al. 2007), making them useful kinematical probes for understanding the structure of galaxies, and to test whether a galaxy contains a substantial amount of dark matter (e.g. Romanowsky et al. 2003; Herrmann & Ciardullo 2009).

There are roughly 3,000 planetary nebulae known (Kohoutek, 2001:1500 PNe, Acker et al., 1992:1143 PNe) in the Milky Way Galaxy and surveys continue to find more.

PN Surveys:

  • Southern Sky:
    • The AAO/UKST Supercosmos H-alpha Survey (SHS):
      UKST Telescope (1.2 m, FOV:6ox6o)
    • The Macquarie/AAO/Strasbourg H-alpha Planetary Galactic Catalog (MASH):
      UKST Telescope (1.2 m, FOV:6ox6o)
      Technique: difference imaging; H-alpha + (:cell PQA(PSS(NII):) 6548, 6584A
  • Northern Sky:
    • The INT/WFC Photometric H-alpha Survey of the Northern Galactic Plane (IPHAS):
      Isaac Newton Telescope (INT, 2.5m)
      Technique: difference imaging; H-alpha, r and i filters

The availability of wide-field imaging surveys of high resolution and sensitivity across di erent wavelength regimes in the optical, MIR and radio in particular, not only provide new discovery media but also give enhanced opportunities to explore the broader multi-wavelength properties of PN and their mimics than hitherto possible (e.g. Parker et al. 2006; Cohen et al. 2007).

HII Regions

An HII region is a large, low-density cloud of partially ionized gas in which star formation has recently taken place.

Much of our current knowledge regarding star-forming patterns and circumstellar disk evolution derives from the study of molecular cloud complexes within a few hundred parsecs of the Sun. Among these are a large number of lower mass clouds such as Taurus and, more infrequently, dense clouds like Orion, which is the prototypical high-mass and high-density star-forming region. While nearby cloud complexes serve as our primary empirical guide to understanding the formation and early evolution of stars, it is important that we study more than just the nearest examples.

The physical appearance of the nebulae on optical images is thought to reflect a combination of a large, background HII regions overlaid by several foreground, dark clouds, with the edges of the dark clouds illuminated by the optically unseen primary exciting star of the HII region.

Fig-2.3-1. (Left) Optical and (Right) near-IR image of some portion of the Orion Nebula.
Image credit: NASA, STScI, ESA.

Updates: 09.05.2012 (Umut)

Star Formation

Stars like our Sun form deep inside molecular clouds. At the very outset of star formation, in order for a star to condense material; the force of gravity (which tends to contract material) must overcome the thermal pressure (which opposely push the interstellar material apart). This material therefore starts to accrete onto a dense center. According to the gas law, the lower pressure regions must be the regions that are as cold as possible. The only regions in the Interstellar Medium to have dense and cold conditions to ignite star formation are the cold and dark molecular clouds. A typical cloud core has a density of n~104 cm-3 and a temperature of T~10K. Once the cloud start to collapse into its own gravity, it begins to produce stellar luminosities, at which point the protostellar phase begins. Star formation is a multi-step process from prestellar core to planet forming disk where it is evolved in time.

With the recent mid-infrared Spitzer and Gould Belt surveys, confirmed number of young stellar objects (YSOs) have increased significantly in the recent years (e.g., Evans et al., 2009; Rebull et al., 2011). Among those, low-mass young stellar objects have traditionally been classified based on their observed infrared (IR) slope, in the wavelength range from 2 to 20 µm (Lada & Wilking, 1984) or their bolometric temperature, Tbol (Myers & Ladd, 1993). In the earliest stages of star formation, most of the emission appears at far-IR and submillimeter wavelengths. In the Class 0 stage (André et al. 2000), which is the earliest and deeply-embedded stage of YSOs, the protostar is surrounded by a collapsing envelope and a circumstellar disk through which material is accreted onto the growing star (Menv >> M). This stage lasts only a short time, typically ~104 – 105 yr, but it is critical for the subsequent evolution since the mass of the star and the physical and chemical structure of the circumstellar disk are determined here. These protostars also power bipolar outflows with extreme characteristics such as a very high degree of collimation and evidence for shock processing of molecular gas even in cases of very low stellar luminosity (Tafalla et al. 2000). In the subsequent Class I phase, lasting a few ~105 yr, the envelope is dispersed by the outflow, the disk grows, the YSO becomes brighter at IR wavelengths and the outflows diminish in power. The radiative transfer of those protostars are mainly determined by the dust in the envelope therefore rich molecular lines are seen in the IR to sub-/mm wavelengths.

In the later phases of star formation a.k.a. T-Tauri phase is one of the groups of pre-main sequence objects where shows protoplanetary disks. Since the surrounding dusty envelope more and more disappeared in this later phase, the central source start to be visible. However, the temperatures are still well below normal stellar temperatures, so most of the radiation is given off in the infrared and yet cannot be observed with the optical instruments. These sources are also showing irregular variablities which creates a lot of interest.

Size also does matter. Comparing with the total lifetimes of a star, the protostellar phase of a low-mass star is (105-7 yr) quite short so that the IR phase is much shorter than the entire lifetime of 109-10 yrs (as seen in optical). However, in high-mass sources, the star spends almost its entire lifetime as cool infrared objects. Specifically, the duration of protostellar phase in high-mass sources are too short comparing with the low-mass counterparts which makes less number of known sample for high-mass sources. Therefore IR surveys are more valid to find and study these objects.

The still valid questions are;

  • What type of clouds or locations within a cloud are best sites for star formation and eventually form a star?
  • What is the star formation efficiency?
  • What is the initial mass function?
  • What is the difference/link between the low- to high-mass star formation?
  • What is the nature of FU Orionis phenomenon? (Hartmann & Kenyon, 1996)


Acker et al., 1992
André et al. 2000, in Protostars & Planets IV
Bekki & Peng 2006
Cohen et al. 2007
Douglas et al. 2007
Evans et al. 2009, ApJS, 181, 321
Hartmann & Kenyon, 1996, AnuRev, 34, 207
Herrmann & Ciardullo 2009
Iben 1995
Kohoutek, 2001
Lada & Wilking, 1984, ApJ, 287, 610
Myers & Ladd 1993, ApJ, 413 L47
Parker et al. 2006
Rebull et al. 2011, ApJS, 196, 4
Romanowsky et al. 2003
Tafalla et al. 2000, A&A, 359, 967

Updates: 13.05.2012 (Selçuk Topal)

Molecular gas and star formation

Almost 40 years ago, interstellar medium (ISM) was thought to be atomic and the molecules were thought as trace abundance. However, after detecting some dominant molecules apart from H2 for the first time, such as CO (Wilson et al. 1970; Penzias et al. 1970; Solomon et al. 1971) the second most abundant molecule in ISM, we realised that the ISM is not only atomic, but also molecular.

How do the molecules form? You cannot simply combine two atoms in the ISM to create a molecule. Because the problem is that it takes more than Hubble time, so more than the age of the universe. So, there must be a 'faster' way to create molecules. The binding energy between the interacting atoms should be carried away by sort of mechanism that molecule can form. The dust grain can do it on their surface. Hydrogen atoms hit the dust, combine with another atom and hydrogen molecules are created. The binding energy is absorbed on the surface of the dust grains. Ones hydrogen molecules form, other more complex molecules can be created relatively easily by hydrogen molecule interacting with other atoms in the ISM. The dust has a crucial role in forming hydrogen molecules and so, the other molecules and finally forming stars and planets that we see today. We now know that ISM hosts atomic, molecular and dusty environment requiring to be probed in UV, visual, IR and radio wavelengths.

Molecules are located deep down the clouds where the density and temperature high enough. Since the light coming from the centre of the cloud will be absorbed by the dust, it is impossible to see behind the dust in optical wavelength. Only the radio emission can pass the dust easily. Molecules cool the clouds by emitting photons, and clouds start to collapse that allows to create necessary physical conditions, such as high gas pressure leading higher temperature, at the core of clouds to form stars. Stars deep inside the cloud heat the gas and dust around, and the dust starts to shine in infrared wavelengths. This is why, not only sub-mm or mm observations, but also infrared observations are complementary tools to get better understanding of star formation occurring in the molecular ISM.

Today infrared observations are used to study on polycyclic hydrocarbon (PAH) molecules in mid-infrared as a star formation rate indicator (Calzetti, 2011), to derive molecular mass through the direct detection of mid-IR H2 emission which is done indirectly in radio using CO emission (Peterson et al. 2012).


Calzetti D., 2011, EAS, 46, 133
Penzias A. A., Jefferts K. B., & Wilson R. W., 1970, AJ, 165, 229
Peterson B. W., et al., 2012, ApJ, 751, 11
Solomon P., Jefferts K. B., Penzias A. A., & Wilson R. W., 1971, AJ, 163, L53
Wilson R. W., Jefferts K. B., & Penzias A. A., 1970, AJ, 161, L43

What can be done with DAG
  • Star formation is still mystery. Specifically T-tauri phase can be very well studied with DAG. Also binary stars has a great expertise in Turkey, however, binary star formation is simply a huge gap in knowledge. Specifically late stages of star-formation is best observed with the IR instruments and combing the expertise with the IR observations of these sources may help understanding the formation.
  • PN surveys requires big FOV (greater than 1 degree) and high pixel scale (less than 1 arcsec).
  • Spectral confirmation of found PNe is needed in optical band with a medium (10000) or high resolution spectrometer (100000 to measure radial velocity in 1 km/s).