star formation

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Star formation is a multi-stage process triggered in the densest cores of giant molecular clouds (with masses up to 10⁶ solar masses). External influences — spiral density waves, shock waves from supernovae, cloud collisions — create local over-densities. Gravitational instability overcomes the combined support of thermal pressure, magnetic fields, and turbulence. The initial isothermal collapse transitions to adiabatic when the core becomes opaque to its own radiation. A protostar forms with an accretion disk and bipolar outflows. When the central temperature reaches millions of degrees, deuterium ignites first, and at ~10⁷ K, hydrogen ignites in the pp-chain or CNO cycle. The star arrives on the main sequence, where it spends almost its entire life. The mass distribution of newborns (the initial mass function) is roughly universal: many low-mass and very few massive stars.

History

Ancient observers considered stars eternal, but already in the 18th century, Immanuel Kant and Pierre-Simon Laplace proposed the nebular hypothesis (from Latin nebula — mist): stars form from rotating gas clouds. In the early 20th century, James Jeans mathematically described the conditions for gravitational collapse. And with the development of infrared astronomy in the second half of the 20th century, astronomers were able to peer into the dense dust and directly observe protostars.

How it works

The main battle in star formation is between gravity, trying to compress the cloud, and thermal pressure, which pushes it outward. If the cloud cools effectively (radiating energy), gravity wins. As it contracts, however, the material spins up due to conservation of angular momentum, so a disk forms around the protostar. Some gas is ejected along the rotation axis in narrow jets. The mass of the newborn star determines its fate: lightweight red dwarfs shine dimly for trillions of years, while heavy giants burn out quickly and explode as supernovae.

💡 The Orion Nebula is the nearest 'stellar nursery' to us (about 1300 light-years away), where among the glowing gas one can see protoplanetary disks like dark silhouettes against a bright background. Planets are assembling in them right now.

The birth of a star begins in a cold (10–30 degrees above absolute zero) and dense molecular cloud, composed mainly of hydrogen and dust. In such clouds, stretching over many light-years, clumps — pockets of slightly higher density — form randomly. Gravity pulls the material together, and it heats up (like a pump compressing air). If the clump can get rid of heat (radiating it as infrared light through dust grains), the collapse continues. At the center, a dense core forms — a protostar, surrounded by a rotating disk. When the core temperature reaches about 10 million degrees, a thermonuclear reaction ignites: hydrogen converts into helium, releasing colossal energy — the star turns on. The remaining disk material can clump together into planets.

History

How it works

💡 In regions like the Orion Nebula, over a thousand stars are born in a few million years, but the efficiency is low: only 10–20% of the cloud's mass turns into stars, the rest is swept away by powerful radiation and winds from the most massive newborns. These same winds, compressing the cloud's edge, can trigger a new wave of star formation.

Star formation is the astrophysical process of transforming interstellar gas into stars. It takes place in the superdense cores of molecular clouds (temperature 10–30 K, density 10³–10⁶ cm⁻³, dominated by H₂ with traces of CO and dust). Key stages: (1) Collapse when the Jeans criterion is satisfied (M > M_J); otherwise the cloud stabilizes. Magnetic flux can slow collapse (critical mass-to-flux ratio). Free-fall time is t_ff ~ (Gρ)^-1/2 ~ 10⁵ years. (2) Formation of the first hydrostatic core (size ~5 AU, ~100 K) and subsequent second collapse upon dissociation of molecular hydrogen, leading to a stellar core. (3) The protostar grows via disk accretion; magnetocentrifugal winds carry away angular momentum. Accretion is episodic (due to gravitational or magnetorotational instabilities). (4) In the pre-main-sequence phase, the protostar contracts along Hayashi or Henyey tracks until hydrogen ignites in the core. Massive stars (≥8 M⊙) reach the main sequence while still accreting. The efficiency of star formation per free-fall time is low (ε ~ 0.01). The initial mass function (IMF) is approximated by a power law dN/dM ∝ M^-α, α ≈ 2.35 for M > 0.5 M⊙, with a log-normal downturn at low masses. Feedback from protostellar jets and radiation regulates efficiency, and massive stars, with their hard radiation and explosions, can disrupt the parent cloud.

History

The nebular hypothesis was proposed in the 18th century by Swedenborg, Kant, and Laplace. In the early 20th century, James Jeans examined the gravitational instability of a homogeneous medium (Jeans criterion). The discovery of dark nebulae by Barnard and others confirmed the existence of cold clouds. Infrared astronomy (IRAS, Spitzer, Herschel) in the 1970s–2010s revealed protostars, and the ALMA interferometer allowed seeing disks and jets in detail. Subrahmanyan Chandrasekhar made a fundamental contribution to the theory of stellar structure and evolution, formulating the Chandrasekhar limit for white dwarfs.

How it works

Star formation is governed by the balance of gravity, pressure, magnetic fields, and turbulence. Jeans criterion: a cloud of mass M collapses if M > M_J = (π/6) c_s³ G^(-3/2) ρ^(-1/2). In magnetized media, the critical mass-to-magnetic flux ratio must be exceeded. Efficient cooling via dust emission in the far-IR ensures isothermality of the early collapse. Once the optical depth reaches about unity, the core heats up. Conservation of angular momentum leads to the formation of an accretion disk; magnetic braking can reduce it. Accretion onto the protostar generates shock waves emitting in the UV and optical. At ~10⁶ K, deuterium ignites; at ~10⁷ K, hydrogen ignites. For massive stars, radiation pressure plays a role, which can halt spherical accretion; material falls through the disk. Bipolar outflows (jets) carry away excess material and angular momentum. On galactic scales, the star formation rate is related to the gas surface density by the empirical Kennicutt-Schmidt law: Σ_SFR ∝ Σ_gas^n, where n ≈ 1.4. Key equations: hydrostatic equilibrium dP/dr = -G M_r ρ/r², continuity equation, energy transport and nuclear energy generation equations, equation of state (ideal gas with radiation pressure), opacity. The initial mass function (IMF) results from turbulent fragmentation and thermal physics; the IMF Salpeter law is often used for high masses.

M_J = \frac{\pi}{6} \frac{c_s^3}{G^{3/2} \rho^{1/2}}

c_s — speed of sound in the cloud (determined by temperature and composition), G — gravitational constant (6.674×10⁻¹¹ m³ kg⁻¹ s⁻²), ρ — gas density

t_{ff} = \sqrt{\frac{3\pi}{32 G \rho}}

G — gravitational constant, ρ — mean density of the cloud

\Sigma_{\text{SFR}} = A \left( \Sigma_{\text{gas}} \right)^n

Σ_SFR — surface density of star formation rate (M⊙ yr⁻¹ kpc⁻²), Σ_gas — surface density of gas (M⊙ pc⁻²), A — efficiency constant, n ≈ 1.4

Open problems & fun facts

What determines the initial mass function and its possible non-universality in extreme conditions?; How does feedback from stars (radiation, winds, supernovae) communicate to the multiphase interstellar medium and regulate the global star formation rate?; What is the role of magnetic fields in the formation of protostellar disks and the launching of outflows?; How do the first stars (Population III) form in metal-free gas, and what are their characteristic masses?; How does the star formation process change in extreme environments — in galactic centers, during mergers, in the early Universe?

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arXiv:2605.13042v1 · 2026-05-13

star formation

History

How it works

History

How it works

History

How it works

Open problems & fun facts

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