Molekuliniai debesys ir protžvaigždės

Molecular clouds and protostars

How cold, dense clouds of gas and dust collapse, forming new stars in stellar nurseries

Between stars, seemingly empty spaces quietly float huge clouds of gas and dust – molecular clouds. These cold, dark regions, embedded in the interstellar medium (ISM), are birthplaces of stars. Gravity can compress matter there enough to trigger nuclear fusion, thus beginning a star's long life journey. From dispersed giant molecular complexes spanning tens of parsecs to compact dense cores – these stellar nurseries are essential to renewing a galaxy's star populations, forming both low-mass red dwarfs and higher-mass protostars that will one day shine brightly as O or B spectral class stars. This article examines the nature of molecular clouds, how they collapse to form protostars, and the subtle interplay of physics – gravity, turbulence, magnetic fields – that governs this fundamental star formation process.

1. Molecular clouds: star-forming nurseries

1.1 Composition and conditions

Molecular clouds are mainly composed of hydrogen molecules (H2), as well as helium and a small amount of heavier elements (C, O, N, etc.). They often appear dark in the visible spectrum because dust particles absorb and scatter starlight. Their typical properties are:

  • Temperature: ~10–20 K in dense regions, low enough for molecules to remain intact.
  • Density: From a few hundred to several million particles per cubic centimeter (e.g., a medium a million times denser than the average interstellar space).
  • Mass: Clouds can range from a few Solar masses to more than 106 M⊙ (in so-called giant molecular clouds, GMC) [1,2].

Such low temperature levels and high densities create conditions for molecules to form and persist, while simultaneously creating a protected environment where gravity can overcome thermal pressure.

1.2 Giant molecular clouds and their subsystems

Giant molecular clouds, extending tens of parsecs, have complex internal structures: filaments, dense clumps, and cores. These subregions often appear gravitationally undefined (can collapse), thus forming protostars or small cluster groups. Observations in the millimeter and submillimeter wavelength range (e.g., ALMA) reveal intricate filamentary structures where star formation often concentrates [3]. Such molecular lines (CO, NH3, HCO+) and dust continuum maps help determine column density, temperature, and motion patterns, showing how subregions can fragment or collapse.

1.3 Factors initiating collapse

Gravity alone is not enough to trigger large-scale cloud collapse. Additional “ignition mechanisms” are:

  1. Supernova shock waves: Expanding supernova remnants can compress the neighboring gas medium.
  2. Expansion of H II regions: Ionizing radiation from massive stars blows shells out of neutral material, pushing them into nearby molecular clouds.
  3. Spiral density wave effect: In galactic disks, passing spiral waves can compress gas, forming giant clouds and later star clusters [4].

Although not all star formation requires external triggering, these processes often accelerate the fragmentation of cloud segments and gravitational collapse in weakly stable regions.

2. The onset of collapse: core formation

2.1 Gravitational instability

If part of the internal mass and density of a molecular cloud exceeds the Jeans mass (the critical mass at which gravity overcomes thermal pressure), that region begins to collapse. The Jeans mass depends on temperature and density:

MJ ∝ (T3/2) / (ρ1/2).

In typical cold, dense cores, thermal or turbulent pressure can no longer withstand gravity, so star formation begins [5].

2.2 The role of turbulence and magnetic fields

Turbulence in molecular clouds promotes chaotic flows that can slow down direct collapse but can also create conditions for local condensations in core locations. Meanwhile, magnetic fields provide additional support if the cloud is pierced by magnetic force lines. Observations (e.g., polarized dust radiation, Zeeman splitting) allow measuring the magnetic field strength. The interaction of gravity, turbulence, and magnetism determines the speed and efficiency with which stars ultimately form [6].

2.3 Fragmentation and clusters

During collapse, the same cloud can fragment into several dense cores. This explains why stars usually form in clusters or groups – the common birth environment can range from a few protostars to rich star clusters with thousands of members. In those clusters, very low-mass brown dwarfs and massive O-type protostars form, essentially born simultaneously in the same GMC.

3. Protostars: formation and evolution

3.1 From dense core to protostar

Initially, a dense core at the cloud center becomes opaque to its own radiation. As it continues to contract due to gravity, heat is released that warms the developing protostar. This object, still embedded in a dusty environment, does not yet perform hydrogen fusion – its luminosity is mainly due to gravitational contraction energy. Observations show that the early protostellar phase is most prominently revealed in the infrared and submillimeter range, since dust obscures the optical spectrum [7].

3.2 Observational classes (0, I, II, III)

Protostars are classified according to the spectral energy distribution (SED) related to dust:

  • Class 0: The earliest stage. The protostar is heavily surrounded by the envelope, accretion is high, and almost no stellar light can penetrate.
  • Class I: The envelope mass is significantly reduced, and a protostellar disk forms.
  • Class II: Commonly called T Tauri (low mass) or Herbig Ae/Be (intermediate mass) stars. They already have prominent disks but less surrounding envelope, and radiation is observed in the visible or near-IR range.
  • Class III: A pre-main-sequence star with almost no disk left. It is close to the final stellar form, with only a faint disk trace remaining.

This classification reflects the star's evolution from a deeply embedded early stage to a more exposed pre-main-sequence star, which will eventually enter the hydrogen fusion phase [8].

3.3 Dipolar outflows and jets

Protostars are characterized by emitting dipolar flows or collimated jets along the rotation axis, which are believed to be caused by magnetohydrodynamic processes in the accretion disk. These flows blow cavities in the surrounding envelope, creating impressive Herbig–Haro (HH) objects. At the same time, slower, wider flows help remove excess angular momentum from the infalling material, preventing the protostar from spinning up too fast.

4. Accretion disks and angular momentum

4.1 Disk formation

As the cloud core collapses, conservation of angular momentum forces the infalling material to concentrate into a rotating circumstellar disk around the protostar. In this gas and dust disk, which can have a radius of tens or hundreds of AU (astronomical units), a protoplanetary disk may eventually form, where planetary accretion occurs.

4.2 Disk Evolution and Accretion Rate

The flow of material from the disk to the protostar is determined by disk viscosity and MHD turbulence (called the “alpha-disk” model). Typical accretion rates can reach 10−6–10−5 M⊙ per year, and as the star approaches its final mass, this rate decreases. By observing the disk's thermal emission in the submillimeter range, astronomers can determine the disk mass and transverse structure, while spectroscopy reveals hot accretion spots on the star's surface.

5. Formation of High-Mass Stars

5.1 Challenges of Massive Protostars

Additional obstacles characterize the formation of high-mass (O and B spectral class) stars:

  • Radiation pressure: The bright luminosity of a protostar creates strong outward radiation pressure, halting accretion.
  • Short Kelvin-Helmholtz timescale: Massive stars heat up their cores very quickly and begin fusion while still accreting material.
  • Cluster environment: Massive stars typically form in dense cluster centers, where interactions, radiation, and jets influence the overall gas evolution [9].

5.2 Competitive Accretion and Feedback

In dense cluster regions, many protostars compete for shared gas resources. Ionizing photons and stellar winds from massive stars can photo-evaporate nearby cores, modifying or even halting their star formation. Despite these challenges, massive stars form—they are the main sources of energy and chemical enrichment in nascent star-forming regions.

6. Star Formation Rate and Efficiency

6.1 Overall Galactic Star Formation

On a galactic scale, star formation (SF) correlates with the gas surface density, as described by the Kennicutt–Schmidt law. Giant star formation complexes can form in spiral arms or bar structures. In dwarf irregular galaxies or low-density regions, star formation occurs more episodically. Meanwhile, in starburst galaxies, short but intense star formation episodes can occur due to interactions or inflows of material [10].

6.2 Star Formation Efficiency

A molecular cloud of gas becomes stars. Observations show that the star formation efficiency (SFE) in a single cloud can range from a few to several tens of percent. Feedback from protostellar outflows, radiation, and supernovae can disperse or heat the remaining gas, halting further collapse. Therefore, star formation is a self-regulating process, rarely converting the entire cloud into stars at once.

7. Protostellar lifetimes and transition to the main sequence

7.1 Time periods

 

  • Protostellar phase: For low-mass protostars, this phase can last several million years until nuclear hydrogen fusion begins in the core.
  • T Tauri / Pre-main sequence: This bright pre-main sequence phase continues until the star stabilizes on the zero-age main sequence (ZAMS).
  • Greater mass: More massive protostars contract even faster and start hydrogen fusion – often within a few hundred thousand years.

 

7.2 Onset of hydrogen fusion

When the core temperature and pressure reach a critical threshold (about 10 million K ~1 solar mass star), hydrogen fusion in the core begins. The star then settles onto the main sequence, where it shines steadily for millions or even billions of years – depending on the star's mass.

8. Current research and future prospects

8.1 High-resolution imaging

Instruments like ALMA, JWST, and large ground-based telescopes (equipped with adaptive optics) allow us to penetrate dusty protostellar "cocoons," revealing disk kinematics, outflow structures, and early fragmentation processes in molecular clouds. As sensitivity and spatial resolution improve, we will gain deeper insight into how small-scale turbulence, magnetic fields, and disk processes interact during star formation.

8.2 Detailed chemistry

Star-forming regions host a complex chemical environment where even complex organic molecules and prebiotic compounds form. Observing the spectral lines of these compounds in the submillimeter and radio ranges allows tracking the evolutionary phases of dense cores – from early collapse stages to protoplanetary disk formation. This relates to how planetary systems acquire their initial volatile resources.

8.3 The importance of large-scale environment

The galactic environment – e.g., shocks caused by spiral arms, bar-driven gas flows, or external compressive factors through galaxy interactions – can systematically alter the star formation rate. Future multi-wavelength observations, combining nearby IR dust maps, CO line fluxes, and star cluster distributions, will better reveal how molecular clouds form and collapse across entire galaxies.

9. Conclusion

The collapse of molecular clouds is a decisive factor in the early stages of a star's life, turning cold, dusty pockets of interstellar material into protostars, which later begin fusion and enrich galaxies with light, heat, and heavy elements. From gravitational instabilities breaking apart giant clouds to the details of disk accretion and protostellar outflows – star formation is a multifaceted, complex process governed by turbulence, magnetic fields, and the surrounding environment.

Whether stars form in isolated environments or dense clusters, the path from core collapse to the main sequence is a universal principle of star formation in the cosmos. Understanding these early phases – from faint Class 0 sources to bright T Tauri or Herbig Ae/Be stages – is a fundamental astrophysical task requiring advanced observations and modeling. A detailed understanding of this interval – from interstellar gas material to a mature star – reveals the key laws that sustain the "liveliness" of galaxies and prepare conditions for planets and possible life in many star systems.

Links and further sources

  1. Blitz, L., & Williams, J. P. (1999). The Origin and Evolution of Molecular Clouds. In Protostars and Planets IV (eds. Mannings, V., Boss, A. P., Russell, S. S.), Univ. of Arizona Press, 3–26.
  2. McKee, C. F., & Ostriker, E. C. (2007). "Theory of Star Formation." Annual Review of Astronomy and Astrophysics, 45, 565–687.
  3. André, P., Di Francesco, J., Ward-Thompson, D., et al. (2014). "From Filamentary Networks to Dense Cores in Molecular Clouds." Protostars and Planets VI, University of Arizona Press, 27–51.
  4. Elmegreen, B. G. (2002). "Star Formation in a Crossing Spiral Wave." The Astrophysical Journal, 577, 206–210.
  5. Jeans, J. H. (1902). "The Stability of a Spherical Nebula." Philosophical Transactions of the Royal Society A, 199, 1–53.
  6. Crutcher, R. M. (2012). "Magnetic Fields in Molecular Clouds." Annual Review of Astronomy and Astrophysics, 50, 29–63.
  7. Shu, F., Adams, F. C., & Lizano, S. (1987). "Star formation in molecular clouds: Observation and theory." Annual Review of Astronomy and Astrophysics, 25, 23–81.
  8. Lada, C. J. (1987). "Star formation – From OB associations to protostars." IAU Symposium, 115, 1–17.
  9. Zinnecker, H., & Yorke, H. W. (2007). "Toward Understanding Massive Star Formation." Annual Review of Astronomy and Astrophysics, 45, 481–563.
  10. Kennicutt, R. C., & Evans, N. J. (2012). "Star Formation in the Milky Way and Nearby Galaxies." Annual Review of Astronomy and Astrophysics, 50, 531–608.
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