IAL 21: Star Formation

Don't Panic

Sections

  1. Introduction to Stellar Evolution as a Prologue to Star Formation
  2. The Stuff in Space
  3. The Interstellar Medium (ISM)
  4. The Intergalactic Medium (IGM)
  5. Molecular Clouds
  6. Star Formation
  7. Disks
  8. The Evolution of Star Formation Regions

  1. Introduction to Stellar Evolution as a Prologue to Star Formation

    2. Stellar evolution is, to say the least, a shaggy dog story: it goes on and on with intricate jargon, parallel stories, and a wide variety of effects. It's hard to keep everything straight---that is yours truly finds it hard to keep everything straight. See the shaggy dog in the figure below (local link / general link: dog_bearded_collie.html).

    To     mix metaphors, we will try to walk a dry path through the morass.

    1. A Quick Overview of Stellar Evolution:

      See the figure above/below (local link / general link: stellar_evolution_overview.html) for an overview of the stellar evolution of stellar evolution of a star of less than ∼ 8 M_☉.

    2. What is Stellar Evolution?

      We must first say that the term stellar evolution is somewhat misleading in that stars do NOT evolve by evolution by natural selection. For evolution by natural selection, see the figure below (local link / general link: evolution_microbes.html).

      Stellar evolution, in fact, applies to the lifetimes of individual stars: their beginnings, courses of development, and final fates: i.e., birth, life, and death. We say they evolve.

      However, there are generations of stars too. Each generation injects into the interstellar medium (ISM) metal-enriched matter out of which new stars form. However, since low-mass stars live a long time relative to high-mass stars, the generations of stars overlap badly.

      Do stars change with generations?

      Yes and no.

      The arena of the generations of       star formation is the observable universe whose evolution is currently mostly adequately described by the Λ-CDM model (AKA concordance model) of cosmology. For the tensions (or anomalies) of the Λ-CDM model, see Tensions of the Λ-CDM Model Since Circa 2018.

      The figure below (local link / general link: cosmos_history.html) illustrates the Λ-CDM model (AKA concordance model) and similar cosmological models.

    3. Observing Stellar Evolution:

      We CANNOT directly observe the whole lifetime (i.e., all the evolution phases) of an individual star.

      We must, in fact, infer stellar evolution theoretical modeling and from observations of stars in different phases of stellar evolution.

      We do have snapshots of most kinds of stars in almost all phases and some rapid changing phases (e.g., explosions of some kind and in particular supernova explosions) can be directly observed.

      Modeling helps to connect the snapshots of stellar evolution in action.

    4. A Longer Overview of Stellar Evolution:

      Because the stellar evolution is a complex, long, and---let's admit it---a difficult-to-remember story, we will start with two cartoons shown in the figure below (local link / general link: diagram/star_life.html) showing the evolution of a low-mass star (the Sun) and a high-mass star (e.g., a 17.5 M_☉ B0 star: Cox-389).

      Nota bene:
      1. Stars less massive than about 8 M_☉ end their lives as white dwarfs; those more massive that about 8 M_☉ end their lives by exploding as supernovae.
      2. Thus, 8 M_☉ is a usual dividing line between low and high-mass stars.

      These cartoons just preview the whole long story.

  2. The Stuff in Space

    3. Space is NOT empty---between the stars and planets that is.

    The density of matter is just very low by comparison to the environment inside stars and planets.

    In this section, we discuss the stuff in Space in general, but mainly in anticipation of discussing the interstellar medium (ISM) (the stuff inside galaxies NOT counting dark matter) which we specialize to explicitly below in section The Interstellar Medium (ISM).

    The interstellar medium (ISM) consists of gas, interstellar dust, electromagnetic radiation (EMR), magnetic fields, and, of course, dark matter.

    Between galaxies is the intergalactic medium (IGM) which we specialize to explicitly below in section The Intergalactic Medium (IGM).

    1. Gas in Space:

      The gas in space is almost always of the cosmic composition: ∼ 73 % hydrogen (H), ∼ 25 % helium (He), and ∼ ≤ 2 % metals (Z)---NOT considering the gas dark matter particles for now.

      This gas is usually NOT obvious in the visible band (fiducial range 0.4--0.7 μm = 4000--7000 Å).

      It is in other wavelength bands: e.g., radio, infrared (IR), X-ray.

      Of course, which wavelength band you see the gas in depends on its thermodynamic state which depends on its temperature, density, impinging electromagnetic radiation (EMR), and other things.

      We will talk about some of these gas components below and elsewhere.

      But we will emphasize important cases of RADIATING gas:

      1. Atomic Hydrogen Gas:

        Dilute atomic hydrogen (H I) makes up most of the interstellar medium (ISM) and typically has a density of ∼ 0.2 to 50 atoms per cm**3 (see Wikipedia: Interstellar medium: Interstellar matter).

        The atomic hydrogen (H I) (when cold as it mostly is) radiates the emission spectral line the hydrogen 21-centimeter line which is the radio band (fiducial range 3 Hz -- 300 GHz = 0.3 THz, 0.1 cm -- 10**5 km). See the figure below (local link / general link: atom_001_h_001_21_cm_line.html) explicating the hydrogen 21-centimeter line.

        Interstellar dust is largely transparent to this emission spectral line, and so it has allowed us to map the Milky Way using radio astronomy starting in the 1950s (see Wikipedia: Hydrogen line: Discovery).

        Hydrogen 21-centimeter line observations are beginning to have an impact on cosmology (e.g., Chowdhury et al. 2020) and is expected to have a great future in the study of the Cosmic Dark Ages (c. 370,000 years -- 150 Myr post-Big Bang) (see Wikipedia: Hydrogen line: In cosmology).

      2. Molecular Hydrogen (H_2):

        Molecular hydrogen (H_2) is 2 hydrogen atom chemical bonded together. It is the common form of hydrogen near the Earth's surface and in human activities. If the hydrogen economy ever becomes a reality, your car might be fueled by molecular hydrogen.

        In space, molecular hydrogen and other interstellar molecules primarily exist in relatively dense environments called molecular clouds where they are shielded by interstellar dust from ultraviolet band (fiducial range 0.01--0.4 μm) which photodissociates them into unbonded atoms.

        Now molecular hydrogen, though overwhelmingly the most abundant molecule in molecular clouds, is nearly invisible: it has virtually NO emission in the thermodynamic state of molecular clouds.

        The other molecules in molecular clouds, though trace amounts by abundance, radiate much more strongly. Their radiation cools the gas which lowers pressure which permits gravitational runaways onto dense cores which become protostars which become stars. It's all a bit like the This Is the House That Jack Built. We elaborate in poetic form on how molecular clouds give rise to star formation in the figure below (local link / general link: star_formation_jack.html).

        We explicate molecular clouds and star formation further in sections Molecular Clouds and Star Formation below.

      3. Carbon Monoxide (CO):

        The near invisiblity of molecular hydrogen means that to observe the inside of molecular clouds---which can only be done in the radio band (fiducial range 3 Hz -- 300 GHz = 0.3 THz, 0.1 cm -- 10**5 km) and infrared band (fiducial range 0.7 μm -- 0.1 cm) because of the opacity of interstellar dust in the ultraviolet band (fiducial range 0.01--0.4 μm) and visible band (fiducial range 0.4--0.7 μm = 4000--7000 Å)---one typically uses spectral lines of tracer molecules: often carbon monoxide (CO) (see Wikipedia: Carbon monoxide: Astronomy). For the carbon monoxide (CO) molecule, see the figure below (local link / general link: carbon_monoxide_c_o.html).

      4. H II regions:

        In most of interstellar space, the gas is relatively cold and low density and NOT in molecular form and emits little electromagnetic radiation (EMR) in the visible band (fiducial range 0.4--0.7 μm).

        The exception is that atomic hydrogen (H I) that emits strongly in the visible and ultraviolet = UV (fiducial range 0.01--0.4 μm) in what are called H II regions.

        Strong UV from hot young blue stars (e.g., OB stars), ionizes atomic hydrogen gas. Ionized hydrogen is called H II in astro jargon.

        Recombination of the electrons and the ionized hydrogen atoms (which are bare protons) gives atomic hydrogens in excited energy levels.

        The decay of the excited energy levels by spontaneous emission causes the emission of EMR.

        The strongest visible emission is the red Hα of the hydrogen Balmer series.

        Thus, H II regions frequently look pinkish in enhanced true color images.

        Since hot young blue stars are usually adjacent to the star forming regions where they formed, H II regions are usually adjacent to star forming regions too.

        The large H II region Tarantula Nebula (AKA 30 Doradus) in the Large Magellanic Cloud (LMC) is shown prominently in the figure below (local link / general link: galaxy_lmc.html).

    2. Cosmic Dust:

      Cosmic dust is the generic name for solid grains of metals (meaning metals in the astro jargon sense) in space.

      There are several kinds: circumplanetary dust (which largely is in the form of dusty planetary rings: e.g., the rings of Jupiter), interplanetary dust, interstellar dust, and intergalactic dust.

      Interstellar dust is the main focus of this IAL 21: Star Formation since it is key ingredient in star formation.

    3. Interstellar Dust:

      The gas in space is usually NOT obvious in the visible, but interstellar dust often is since it is very opaque.

      Interstellar dust can often be seen as prominent dust lanes in images of galaxies---see the Sombrero Galaxy (M104, NGC 4594) in the figure below (local link / general link: galaxy_sombrero.html).

    4. What Is Interstellar Dust?

      Now what is that darn interstellar dust?

      Dust grains are believed to form in ejecta from stars in the form of stellar winds and supernovae.???

      This ejecta is rich in refractories (i.e., materials, including metals in the astrophysical context, that condense at relatively high temperatures).

      The refractories condense to form dust grain cores.

      Surrounding the refractory grain cores, volatiles (i.e., materials that condense at relatively low temperatures) can condense to complete a grain consisting of a core and outer layer of different composition.

      Such condensation of volatiles happens primarily in the inner regions molecular clouds where the outer regions of interstellar dust can protect the inner regions from ultraviolet light (fiducial band range 0.01--0.4 μm) which would evaporate the volatiles.????

        The terms refractory and volatile can be used in a relative sense.

        But conventionally in astrophysics certain materials are usually considered to refractories and certain others volatiles.

        But iron, silicate (silicon oxygen substances which make up most ordinary rock), and carbon are ordinarily refractories in our astrophysical contexts.

        Hydrogen, helium gas, other noble elements, N_2, carbon dioxide, water (H_2O), methane (CH_4), ammonia (NH_3) are examples of volatiles in our astrophysical contexts.

        See Se-418.

      The density of metals is high in the aforesaid ejecta and this is what allows the dust grains to grow by condensation---one atom or molecule at a time.

      How much interstellar dust is there compared to interstellar gas?

      Well, interstellar dust is made of metals---except for some hydrogen bonded in molecules. Also stars (which from from the ISM) are at most ∼ 3 % ??? metals (see David Weinberg 2016, "On the Deuterium-to-Hydrogen Ratio of the Interstellar Medium", p. 3, but this is NOT a best reference). So the ISM can only be a few percent metals by mass fraction, and so the interstellar dust can only be a few percent of the ISM usually.

      We do have examples of interstellar dust in our hands to examine. They are called presolar grains (see also Wikipedia: Cosmic dust Stardust) and they are found in primitive meteorites (mostly??? chondrites some of which survive today as the most chemically unprocessed left over from Solar System formation (4.6 Gyr BP)). The properties of presolar grains turn out to be very heterogeneous since the presolar grains form from the ejecta of particular post-main-sequence stars supernovae. What we see of interstellar dust is an average of heterogeneous properties and this average varies depending on the context: i.e., the type of galaxy, region of galaxy, and cosmic time.

      As well as presolar grains, certain Solar System particles may resemble interstellar dust (HI-275).

      See the figure below (local link / general link: interplanetary_dust.html) for an image of a Solar System interplanetary dust particle.

      We do have models of       interstellar dust that result from a combination of remote observations of interstellar dust, laboratory experiments, knowledge about elemental abundances and the common elemental compounds, and theory (Ze2002-316).

      Note silicates (or more properly silicate minerals) are compounds containing silicon and oxygen are what most ordinary rock is made of.

    5. Seeing Through Instellar Dust:

      Although interstellar dust is highly opaque to visible light and ultraviolet, we can only see into and through molecular clouds in the infrared, microwave, and radio.

      A famous dusty nebula is shown in the figure below.

        ./formation/noao_horsehead_nebula_003.jpg

        Caption: The Horsehead Nebula NOAO, Kitt Peak, Arizona.

        The Horsehead Nebula because of its accidental resemblance to a horse's head is one of the most famous of all nebulae.

        This is approximately a true color image. The Horsehead Nebula is a dark dusty nebula. The difference in star counts above and below the mid-image line illustrates that no or few stars are seen through the dark nebula: most of the stars in the lower image are probably foreground.

        The pink color is an emission nebula. Ultraviolet light from hot young stars (mainly OB stars) creates excited hydrogen atoms which then emit strongly in the hydrogen red line (i.e., the H alpha line).

          The excitation is usually accomplished by ionization of hydrogen followed by recombination to an excited level.

        Bright red dots at the base of the Horsehead are protostars. The Horsehead Nebula is also a star formation region.

        Credit/Permission: © N.A.Sharp/NOAO/AURA/NSF, 1994 / NOAO/AURA Image Library Conditions of Use.
        Download site: NOAO: The Horsehead Nebula.
        Image link: Itself.

    6. A Digression on Allotropes of Carbon:

      Graphite is the common form of pure carbon (i.e., allotrope of carbon) used in pencil leads (i.e., graphite cores).

    7. Electromagnetic Radiation in Space:

      Additionally, space is pervaded by electromagnetic radiation (EMR) in all wavelength bands: sometimes a very low amount.

      Between the obvious bright objects---stars and galaxies, it is called the diffuse extragalactic background radiation (DEBRA).

      The figure below (local link / general link: diffuse_extragalactic_background_radiation.html) explicates DEBRA.

      The largest component of       DEBRA is the cosmic microwave background (CMB) which is explicated in the figure below (local link / general link: cmb.html).

    8. Magnetic Fields in Space:

      Then there are magnetic fields---which yours truly tends to call B fields since the conventional symbol for magnetic field is B.

      Magnetic fields are everywhere in space and concentrated ones occur in many important cases: around some planets (like the Earth), associated with neutron stars and black holes, associated with active galaxies and supermassive black holes, and associated with star formation.

      Increasingly magnetic fields are being recognized as important factors in many astrophysical environments.

      Which is a vast complication since the effects of magnetic fields and how they evolve are very complex.

      Recall, we will skirt magnetic fields in IAL as much as possible. They are mainly just beyond our scope.

    9. Dark Matter, Baryonic Dark Matter, and Dark Energy:

      In fact, the stellar matter is only about 0.4 % of the total mass-energy in the observable universe.

      Dark matter and baryonic dark matter are much more abundant than stellar matter.

      So far the dark matter is only known through its gravitational effect on galaxies and the large-scale structure of the universe.

      We discuss dark matter in IAL 27: The Milky Way, IAL 28: Galaxies, and IAL 30: Cosmology.

      Here we will just say that dark matter is probably some kind of exotic particle that is very unreactive, except for gravity. There are theories as to what this exotic particle is, but NO established theory. We hope to discover what it is in the laboratory someday.

      The figure below (local link / general link: pie_chart_cosmic_energy.html) displays the estimated distribution of mass-energy among the various major forms of mass-energy in the observable universe.

  3. The Interstellar Medium (ISM)

    4. The space matter inside galaxies is called the interstellar medium (ISM)---it is, in fact, very complex. We explicate the ISM in the insert below (local link / general link: interstellar_medium_ism_table.html).

    Before leaving off with the     interstellar medium, a question:

  4. The Intergalactic Medium (IGM)

    5. Although it is NOT directly part of star formation, the intergalactic medium (IGM) is the ultimate source of the baryonic matter that goes into star formation. So this is a reasonable point for an explication of the IGM. For the explication, see the figure below (local link / general link: intergalactic_medium.html).

  5. Molecular Clouds

    6. The component of interest of the ISM for star formation is the molecular cloud component.

    Note clouds of all kinds in space are often called nebulae which is just Latin for clouds.

    Objects that appear cloud-like were historically called nebulae too.

    Spiral galaxies, for example, (before some stars in them were resolved) were once called spiral nebulae and still can be if one is speaking historically.

    1. Molecular Clouds:

      Molecular clouds are irregular and turbulent with all kinds of motions and rotations. In some ways they are like sky clouds, but in other ways NOT, of course.

      Molecular clouds are also cold and dark.

      They are dark, because they have heavy concentrations of interstellar dust which makes them opaque in the visible and ultraviolet.

      There's nice little molecular cloud in the figure below (local link / general link: molecular_cloud_rude.html).

    2. The Formation of Molecular Clouds:

      Dense clouds of interstellar dust (which usually are associated with dense concentrations of interstellar gas) are necessary for the formation of molecular clouds, and so for star formation. To explicate:

      1. The dust absorbs visible light and ultraviolet light and radiates it away as infrared light. This shields the interior of the dense clouds of interstellar dust from said visible light and ultraviolet light that would act strongly to break up molecules (CM-298).

      2. The interstellar dust may also act as a catalyst in molecule formation: atoms that stick to the dust may have an easier time combining into molecules which then escape back into space (CM-298; Ze-337).

      3. The molecules formed in dense clouds of interstellar dust make these clouds molecular clouds.

      4. The molecules in turn are efficient radiators of heat energy in the infrared 0.7 μm -- 0.1 cm, the microwave 0.1--100 cm, and the radio 0.1 cm -- 10**5 km to which the interstellar dust is fairly non-absorbing. The cooling provided by the molecules reduces the gas pressure, and so allows fragments of molecular clouds to collapse under self-gravity to dense cores which continue to collapse to form stars (HI-375; CM-119).

      5. Note that although molecular hydrogen (H_2) is overwhelmingly the dominant molecule in molecular clouds, it is a poor radiator. So it is other molecules with trace abundance like carbon monoxide (CO) that are the main coolants????.

      Over 100 kinds of molecules have, in fact, been observed in interstellar space, mostly in molecular clouds??? (HI-373; Wikipedia: List of molecules in interstellar space).

      In fact, organic molecules (which just means they contain carbon, NOT that they a biotic origin) have been observed in molecular clouds and this is suggestive that such clouds may seed planetary systems with building blocks for life??? (HI-373).

      There has even been speculation that molecular clouds could harbor life. Fred Hoyle (1915--2000), well known both as an astrophysicist---he coined the term Big Bang, but in mocking the theory---and a science fiction writer, wrote a noted scifi book The Black Cloud (1957) about an intelligent molecular cloud entity---NOT a very likely story, I'd guess.

    3. Where are Molecular Clouds?

      Almostly only in spiral galaxies and irregular galaxies.

      In spiral galaxies, the molecular clouds are almost only in the spiral arms which are grand design spiral arms (which are spiral density waves) or flocculent spiral arms (which are just wound-up giant molecular clouds (see IAL 28: Galaxies: Spiral Arms).

      In irregular galaxies, the molecular clouds are rather randomly located.

      See the three figures below for molecular clouds and star forming regions in the Milky Way (local link / general link: milky_way_map.html), the irregular galaxy NGC 1427A (local link / general link: ngc_1427a_irr.html), and in galaxies of the Hickson Compact Group (local link / general link: galaxy_hcg_87.html).

    4. The Composition of Molecular Clouds:

      Now the composition of molecular clouds should be approximately the mean cosmic composition which is approximately the same as that of the Sun.

      See the figure below (local link / general link: solar_composition.html) for the primordial solar nebula composition or, for short, the solar composition.

        Question: What is the main molecule in molecular clouds?

        1. molecular hydrogen (H_2).
        2. molecular helium (He_2).
        3. carbon monoxide (CO).











        Answer 1 is right.

        If the most abundant species can form a molecule, it seems reasonable that it forms the dominant molecule.

        By the way, H_2 is the common terrestrial form of hydrogen and the one that one day you may be pumping into your hydrogen-powered car (Wikipedia: Hydrogen economy).

        Helium is a noble gas. It doesn't combine in any kind of molecule, except perhaps in some exceptional cases which we won't go into. Helium in gaseous form always consists of free helium atoms---except perhaps in some exceptional cases which we won't go into. It's a monatomic gas.

        Carbon monoxide (CO) is just a trace gas, but it is observationally important as we discuss just below.

      Unfortunately, H_2 doesn't absorb or emit much EMR in the infrared 0.7 μm -- 0.1 cm, the microwave 0.1--100 cm, and the radio 0.1 cm -- 10**5 km (Ze-331), where molecular clouds are strong emitters because of their low temperatures (of order 10 to 100 K) and where they are very non-absorbing.

      Thus, H_2 is practically invisible in molecular clouds.

      The helium gas is mostly invisible too.

      But many other radio emitting molecules exist in minute amounts in molecular clouds.

      Carbon monoxide (CO) is readly observed and acts as a tracer of the density (CK-300).

      The H_2 and He gases must be inferred from modeling and knowledge of the average cosmic composition.

    5. Giant Molecular Clouds:

      The largest molecular clouds are huge: these are giant molecular clouds (GMCs):

      1. 15 to 60 pc in size scale (see Se-212).
      2. 10**2 to 10**6 M_☉ in mass (see Se-212).
      3. lifetimes of order 10--30 Myr (see Kruijssen et al. 2019).
      4. star formation efficiency ∼ 2--3 % in giant molecular clouds (see Kruijssen et al. 2019).

      Note the nubmers above are NOT definitive. The numbers tend to vary with reference.

      Many molecular clouds are much smaller than giant molecular clouds (GMCs).

  6. Star Formation

    7. In molecular clouds, there can be RUNAWAY GRAVITATIONAL COLLAPSES that lead eventually to star formation.

    The time scale for formation is tens of millions of years, and so we DO NOT observe star formation happening in time.

    But there are many sites of star formation in various stages, and so from these SNAPSHOTS and modeling we can understand the time evolution star formation to a degree.

    Tens of millions of years is also the time scale for lifetime of molecular clouds before they are dispersed or evaporated largely by feedback from star formation itself as we discuss below in section The Evolution of Star Formation Regions.

    The star formation efficiency = a few percent: i.e., the amount of mass of a molecular cloud that goes into new stars in the molecular cloud lifetime.

    Now how do stars form?

    Molecular clouds have a tendency to collapse under their own self-gravity. But that tendency is resisted by various things (HI-334--336; Se-221):

    1. The cloud gas pressure.

      The interstellar dust is necessary to reducing this pressure (see the above subsection The Formation of Molecular Clouds).

    2. Magnetic field pressure. A subject we will NOT go into, but it is somewhat important.

    3. Rotational kinetic energy of the clouds. There is a "sideways motion" that causes objects to keep missing the central mass trying to pull them to it.

      In physics, this "sideways motion" is quantified as angular momentum.

      Recall rotational kinetic energy keeps the planets from falling into the Sun. The planets have only weak means of dissipating their rotational kinetic energy, and so revolve quasi-perpetually.

      But viscosity (from particles rubbing on particles) in molecular clouds helps to reduce the angular momentum effectively and this leads to inspiral. Magnetic fields can help with inspiral???.

    Some TRIGGERING EVENT is needed it is thought to fragment the cloud and increase the DENSITY of a fragment to the point of a RUNAWAY GRAVITATIONAL COLLAPSE.

    TRIGGERING EVENTS are illustrated in the cartoon in the figure below (local link / general link: star_001_collapse.html).

    In a RUNAWAY GRAVITATIONAL COLLAPSE:

    1. the mass compacts into a smaller region,
    2. thus its self-gravity is higher because gravity increases as distances decrease,
    3. this causes more compaction,
    4. this causes higher self-gravity,
    5. and so on.
    6. energy gets radiated away in some form which reduces pressure necessary to continue the collapse.

    Inside the dense regions of runaway collapse, dense cores that fragment to become protostars. See the figure below.

    We CANNOT see the dense cores in the visible band (fiducial range 0.4--0.7 μm) because they are hidden by the interstellar dust.

    But the dense cores heat up and radiate EMR approximately like a blackbody in the radio to which the interstellar dust is relatively transparent.

    The cartoon in the figure below illustrates the energy transformations that occur in making a protostar.

    The observed protostars are of order a solar mass and extend over of order a third of a light-year (Se-220).

    The term protostar is usually used only when the collapsing fragment of a dense core gets sufficiently hot to start radiating in the infrared (IR) (Se-222).

    The protostar phase ends when the star starts the nuclear burning of hydrogen to helium (Se-222).

    Like dense cores, the protostar are hidden by gas and dust from visible observations.

    Even though rather cool, protostars are very luminous because of their large radiating surface and large amount of heat energy that comes from the transformation of gravitational potential energy on collapse.

    Protostars that will become low-mass main-sequence stars are MORE luminous than those stars.

    Protostars that will become high-mass main-sequence stars are typically about as luminous than those stars.

    The pre-main-sequence phase of stars like all phases decreases in time period as mass increases.

    The cartoon Hertzsprung-Russell (HR) diagram in the figure below illustrates the features just discussed. It is based on model calculations, of course.

    A protostar continues to grow in mass, contract in size, and heat up by accretion, but eventually it "bites the hand that feeds its it".

    It develops a strong stellar wind probably owing partially to radiation pressure and partially to magnetic field effects (????), and blows away most of its cocoon of dust and gas (Se-223).

    The outward flow is often BIPOLAR due in part to a protoplanetary disk (see section Disks below) channelling flow out along the poles which are approximately perpendicular to the plane of the protoplanetary disk. These poles are probably approximately aligned with the rotation axis of the protostar.

    The BIPOLAR FLOW is probably affected by magnetic effects.

    The protostar still contracts and gets hotter and at some point, its core will start the nuclear burning hydrogen to helium---at which time one can stop calling it a protostar (Se-222).

    After a bit of settling down evolution, the object becomes a main-sequence star: i.e., a stable, relatively unchanging, hydrogen burning star.

    The time for evolution from the initial stage of a protostar to main-sequence star is strongly mass-dependent as the Hertzsprung-Russell (HR) diagram in the figure above indicated and as given below: 

      A 30 solar mass star takes about 30,000 years.

  A solar mass star takes about 30 million years.

  A 0.2 solar mass star takes about a gigayear. (Is this right??? Seems a bit long.)
    Reference Se-223. These numbers are subject to revision.

  7. Disks

    8. Protostars seem to be surrounded by disks almost always: protoplanetary disks.

    Now to explicate protoplanetary disks.

    1. Images of Protoplanetary Disks:

      Protoplanetary disks have been observed since the 1980s. See the figure below of Beta Pictoris (AKA Betapic) and its famous protoplanetary disk.

      Before that they were entirely theoretical, but they had seemed a reasonable idea since Immanuel Kant (1724--1804) first proposed that planets formed out of disks in 1755 (No-406) based on qualitative reasoning from Newtonian physics. See below subsection From Protoplanetary Disks to Planets.

      A more detailed picture of a protoplanetary disk is shown in the artist's conception of one protoplanetary disk in the figure below.

      Real images of protoplanetary disks confirm to a large degree the artist's conception of one protoplanetary disk in the figure above.

      For real images of protoplanetary disks, see the images taken by ALMA in the figure below (local link / general link: protoplanetary_disks_alma.html).

    2. The Formation of Protoplanetary Disks:

      Protoplanetary disk formation by a relaxation process (plus ancillary processes) is explicated in the figure below (local link / general link: protoplanetary_disk_formation.html).

    3. Accretion Disks in General:

      Protoplanetary disks are a subset of astrophysical accretion disks. Accretion disks are natural in infalling or tidally disrupted astrophysical contexts and their formation process is the relaxation process (plus ancillary processes) as outlined in the figure above (local link / general link: protoplanetary_disk_formation.html), mutatis mutandis.

      Of course, accretion disks happen when some balance of angular momentum, rotational kinetic energy, cooling, and self-gravity do NOT lead to collapse to or onto spherically symmetric object.

      The following are examples of accretion disks:

      1. Protoplanetary disks which are the main topic of this section.

      2. The galactic disks of spiral galaxies and lenticular galaxies

        In the collapse of a large gas cloud to a spiral galaxy (or lenticular galaxy if they form directly), an accretion disk forms which becomes the star and gas/dust disk of the spiral galaxy. Most of the stars form after collapse.

        The spiral arms of spiral galaxies further explanation.

        Elliptical galaxies are different in that star formation happens too rapidly for the cloud of gas and dust to relax to a disk or, probably more usually the galactic disks are destroyed by galaxy mergers that form elliptical galaxies. In the first case, gravitational interactions among stars do NOT dissipate kinetic energy to heat much, and so do NOT allow for a relaxation to a disk.

        We discuss spiral galaxies and elliptical galaxies in IAL 28: Galaxies.

      3. Accretion disks around black holes where the accretion disks form from matter accreted from a close binary companion as discussed in IAL 25: Black Holes.

      4. Planetary rings (e.g., Saturn's rings) can also be regarded as accretion disks though it may NOT be usual to do so.

        Planetary rings are discussed briefly in the figure above (local link / general link: protoplanetary_disk_formation.html) and in IAL 15: Gas Giants. But for a preview of planetary rings, see the figure below of Saturn (local link / general link: saturn_rings_orientation_perspective.html).

    4. From Protoplanetary Disks to Planets:

      The protoplanetary disks around protostars eventually can partially coalesce into planets, but that is the story of IAL 10: Solar System Formation.

      They may NOT always do so, but we now know that planet formation is pretty common. See IAL 18: Exoplanets & General Planetary Systems.

      The figure below (local link / general link: immanuel_kant.html) shows one of the pioneers of the theory of planet formation in protoplanetary disks.

  8. The Evolution of Star Formation Regions

    9. Star formation regions have complicated evolutions:

    1. Some relatively random history brings about a concentration of gas and dust. All of astrophysical history on the smaller scales has a large element of randomness in it. The concentration can become a molecular cloud.

      But in spiral galaxies the relatively random formation of molecular clouds is almost always in the spiral arms. In irregular galaxies, the formation of molecular clouds can be anywhere.

    2. The molecular clouds can become a star formation region. Such regions are rather turbulent and chaotic. They come in all kinds of shapes and sizes, and so no two will ever be exactly the same. See again that nice little molecular cloud in the figure below (local link / general link: molecular_cloud_rude.html).

    3. Some massive, hot, bright, blue OB stars probably form first and they provide strong ultraviolet (UV) light that tends to heat up the gas of dark nebula and evaporate the dust. This tends to turn off further star formation in the immediate vicinity of the OB stars by increasing the pressure that prevents collapses.

      Also the direct pressure of the radiation and stellar winds??? from the OB stars tends to blow away the local material that could go into forming stars.

      If there are a large number of OB stars relatively close together, they could form an OB association.

      The UV light from the OB stars causes hydrogen to become ionized and on recombination in an excited state, the hydrogen emits a hydrogen line spectrum which in the visible consists of the hydrogen Balmer lines. The light in the visible is often red or pink because the strongest hydrogen Balmer line is the red Halpha.

      Because ionized hydrogen is conventionally called H II in astronomy, such emission regions are called H II regions.

    4. Although in the immediate locality the pressure increase caused by OB stars tends to turn off star formation, the pressure of radiation and winds from OB stars can initiate star formation nearby.

    5. OB stars have short main-sequence lives and in under about 10 million years (Se-247??) some of them will blow up as supernovae which create expanding ejecta that can both disperse star forming clouds and compress them to start new star formation.

      Actually, supernovae may only play a secondary role in controlling the behavior of star forming regions (see Kruijssen et al. 2019). As short as time to supernova explosion is from formation for OB stars, the time scale for OB stars to turn off and disperse star forming regions by stellar winds and radiation is shorter it seems: of order 1.5 million years. However, supernovae are still very important in providing new metals to the interstellar medium (ISM) and in driving outflows from galaxies to the intergalactic medium (IGM) or intracluster medium.

    6. The relatively non-local initiation of new star formation by successive short-lived generations of OB stars either in their lives or, at least sometimes, as supernovae (though Kruijssen et al. 2019 may disagree) can cause self-propagation of star formation through a large molecular cloud.

    7. After probably of order 10 million years or so of existence (HI-376; Kruijssen et al. 2019), a star formation region will end up being partially used up and partially dispersed. The fragments of the cloud can end up becoming parts of new star formation regions.

      The efficiency of star formation (i.e., the fraction of molecular cloud that turns into stars before the molecular cloud disperses) is rather low. A fiducial star formation efficiency is 1 % ????, but the actual average efficiency may be more like 2 to 3 % with a range of variation from 0.2 to 20 % or more??? (e.g., Murray 2010; Kruijssen et al. 2019; Kretschmer & Teyssier 2019).

    The complex feedback mechanisms of the OB stars and supernovae make star formation an immensely complicated process. Additionally, the history of a star formation region is somewhat random.

    As the above discussion indicates, dark molecular clouds with star formation are often in proximity to bright emission clouds (H II regions) powered by OB stars that formed out of the clouds at an earlier time. The result is a bit of a mess.

    A famous image of the mess of star formation is the one of Eagle Nebula (M16) with the Pillars of Creation taken by Hubble Space Telescope (HST, 1990--c.2040). See that image and complementary images including one from James Webb Space Telescope (JWST, 2021--2041?) in the figure below (local link / general link: eagle_nebula_large_noao.html).

    OB stars will blow up as supernovae leaving a neutron star or black hole remnant if their mass is greater than about 8 M_☉. Those OB stars less massive than 8 M_☉ (which are only B stars actually), will end their lives as white dwarfs. In both cases, they all their nuclear burning lives (of order 10 million years or less) in or near the star formation region where they were born.

    The compact remnants (white dwarfs, neutron stars, and black holes) of hot young stars (mainly OB stars) will, of course, wander from away from their birth star formation regions which will disperse after tens of millions of years or so (HI-376) as mentioned above.

    Less massive stars like the Sun can live billions of years. They will NOT explode as a supernova and will wander far from their place of formation and often from their sibling stars. Their star formation regions are likely long dispersed by the time such stars are middle-aged.

    Thus we can know little about the individual histories of long-ago formed stars. Also their orbital evolution since formation has been somewhat chaotic.

    We can only have a general account of the average star formation history for such stars.

    This is, of course, even true of Sun and Solar System about which we have so much information.