FROM MOLECULAR CLOUDS TO STARS PAOLO SARACENO AND RENATO ORFEI ISTITUTO DI FISICA DELLO SPAZIO INTERPLANETARIO, CNR Area di Ricerca di Roma Tor Vergata Via del Fosso del Cavaliere 100 – I 00133 Roma Abstract In this short review we outline some aspects of the present understanding of the star formation process, discussing the different processes that matter undergoes from the diffuse clouds to the pre-Main-Sequence stars. We will discuss both the global properties of the star formation process and the evolution of protostellar objects. 1 Star formation and the evolution of matter in the Universe The understanding of the star formation process is critical in astrophysics because it is necessary for the comprehension of the evolution of stars, galaxies and finally of the Universe. Star formation is crucial for the energetics and chemistry of the interstellar medium, that may have strong implications on the origin and evolution of life. Eventually, life requires planets for its existence and the formation of planetary systems can only be studied as a part of the star forming process. The understanding of the star formation process will make clear the origin of the mass spectrum of stars, a crucial point of the evolution of the matter in the Universe. The cycle that matter follows, since the origin of the Universe, is schematically represented in fig. 1. Stars, once formed, are producing, through nuclear reactions, elements heavier than Hydrogen. At the end of their life, part of the stellar mass will return to the interstellar medium (right side of the figure) to form new stars, part of it will be lost and will remain blocked in the remnants (White Dwarfs, Neutron Stars, Black Holes; left part of fig. 1). Fig. 1 shows the importance of mass in the evolution of stars, in fact stars of less than ~0.01 M¤ (Dantona & Mazzitelli 1994) will never reach the temperature necessary to start nuclear reactions, and they will undergo a continuous contraction (Brown Dwarfs). More massive stars will produce energy through nucleosynthesis, their luminosity is a function of the mass as L ∝ M4, therefore the star life is ∝ M-3. This has a consequence for the search of extraterrestrial life: only stars of relatively low mass (less than ~ 4 M¤) may host life in their planetary Proceedings of “The bridge between the bang and biology” Stromboli – September 13 – 17, 1999systems. Life needs a few billion years to develop and massive stars do not last so long. The final stages of stars also depend on stellar masses: stars of less than ~7 Mo will end their life as White Dwarfs. Among them, those of more than ~0.7 M¤ will eject, during the post Main Sequence phase, up to 80% of their initial mass (the fraction increases with mass), passing through the stages of Red Giant and Planetary Nebulae. Star more massive than ~ 7 M¤ will explode as supernovae Figure 1 shows schematically, right part of the figure, the cycle that matter follows since the origin of the universe, enriching at each cycle the interstellarmedium of heavier elements. Ms indicate the original mass of the star and Mc the residual mass after the mass lossproducing a large amount of heavy elements (all the atoms heavier than Fe can only be produced by Supernovae). The explosion ejects the largest part of the mass of the stars in the interstellar medium leaving a relatively small fraction (<10 %) of the initial mass in Black Holes or Neutron Stars. In conclusion: the less massive stars last for a longer time, but at the end of their life most of their masses is lost to the cycle. Increasing their masses, stars will last less (life ∝ M¤–3), but they increase the fraction of matter returned to the interstellar medium, enriched of heavy elements, to form new generations of stars. Therefore the chemical evolution of the Universe is mainly due to massive stars. The evolution of the matter ejected by stars (see Evans 1999 for a review) is schematically represented in the right side of fig. 1 (and will be discussed in the following sections). The matter aggregates in Giant Molecular Clouds, which show density enhancements in relatively small, cold regions called ”Dense Cores”. Stars form inside these cores and the mass of the dense cores settles the upper limit of the star masses. The masses of the Dense Cores have a lower limit given by the Jeans Mass (see § 3), that depends on the temperature of the Cores which depends on the cooling capacity of the gas. The cooling capacity of the gas of the Dense Cores depends on the abundance of molecules, like CO, with a dipole momentum that can be excited at very low temperatures (H2, the most abundant molecule, has only quadrupole momentum). The abundance of these molecules depends on the heavy elements produced in the previous generation of stars. So the mass spectrum influences the production of heavy elements and the production of heavy elements influences the mass spectrum. In the Early Universe, when heavy elements were not abundant, the gas of the clouds should have been much warmer; the Jeans Mass (§ 3) much higher; therefore stars much more massive than those we have nowadays should have been produced (Stahler, 1986). These first generations of massive stars have rapidly increased the abundance of the heavy elements of the Early Universe. 2 Giant Molecular Clouds Giant Molecular Clouds (GMC) are large condensations of gas and dust which are detected and mapped at 2.6 mm by the J = 1-0 rotational transition of the CO molecule (Blitz & Williams 1999; Myers 1999). The average properties of GMCs are: Mass M 104 - 106 M¤ Size 10 - 100 pc Density N ~100 cm-3 Temperature T ~ 10 K Sound speed c 0.2 km/s Gas/dust (in mass) 100 Magnetic field ≤ 10 µGFigure 2: (from Blaauw 1991). The CO map of the Orion cloud (Maddalena et al. 1986) with the subgroup of OB associations formed sequentially. The major component of GMCs is H2, but they are detected in the CO line because H2 (as told above) can not be excited at the low temperature of these clouds. CO is the best tracer of the diffuse medium because it is very abundant ([CO]/[ H2], ~ 10-4), it has a relatively high dissociation energy (11.09 eV) and it can be excited from millimetre wavelengths (few Kelvin) to the near infrared (~ 3000 K) in a variety of different physical environments. In fig. 2 (from Blaauw 1991), the CO J = 1-0 map of the Orion complex (Maddalena
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