Wikipedia

User:ZaperaWiki44/sandbox/Quasi-star

< User:ZaperaWiki44 | sandbox
"Black hole star" redirects here. For black holes created from stars, see Stellar black hole. For stars that become black holes, see Supernova. For types of stars denser than neutron stars, see Exotic star. For the 1994 Soundgarden song, see Black Hole Sun.
Not to be confused with a quasi-stellar object or quasar.
Size comparison of a low-mass quasi-star to some of the largest known modern stars.

A quasi-star[1] or quasistar[2] (QS),[1] also called a black hole star,[3] is a hypothetical type of extremely massive and luminous star that may have existed early in the history of the Universe. Unlike modern stars, which are powered by nuclear fusion in their cores, a quasi-star's energy would come from material falling into a black hole at its core.[4] The formation of such object would have resulted from the core of a large supermassive protostar collapsing into a stellar-mass black hole, where the outer layers of the protostar are massive enough to absorb the resulting burst of energy without being blown away or falling into the black hole, as occurs with supernovae.

Quasi-stars were first proposed in 2006;[2] although a confirmed observation has not yet been made, potential sightings of these objects have been reported by the James Webb Space Telescope since its launch.[5] The study of quasi-stars would provide valuable insight into the early universe, galaxy formation, and the behavior of black holes, namely because they are considered as possible progenitors of the modern supermassive black holes that formed soon after the Big Bang, such as the one in the center of the Galaxy.[4][1][verification needed]

  • This structure is dubbed a quasistar, since it resembles a (scaled-up) red giant in structure, though it is powered by accretion into a central BH (Begelman et al. 2006, 2008, hereafter BRA08).[6]
    • The resulting object, which we have dubbed a ‘quasi-star’ (QS), resembles a red giant with a luminosity comparable to a Seyfert nucleus.[7]

Structure

edit
 
Extreme mass loss usually lead to an extensive and complex circumstellar environment of gas and dust.
 
Hubble Space Telescope image showing the bipolar Homunculus Nebula which surrounds Eta Carinae

Although different models exist for the structure of a quasi-star, it is generally thought to be comprised of a photospheric star-like envelope (that contains the bulk of the star's mass) surrounded by a radiative layer and a core, which includes a black hole embryo within.[6][8][2] This overlying envelope is highly convective and supported by radiation pressure. According to recent models, it may also become adiabatic during a quasi-star's late stage,[6][9] with an inner saturated-convection region forming by that time, thereby conforming to a convection-dominated accretion flow around the black hole.[9] The quasi-star would also have a porous atmosphere around the convective envelope, where convection becomes inefficient.[6] In a 2011 model for the structure of a super-Eddington quasi-star with masses over a million M☉ where it would suffer substantial mass loss due to its extreme luminosity, its strong stellar wind appears as optically thick surface and hence appears as a cooler extended pseudo-photosphere above the porous atmosphere, where the optical density of the wind drops to near zero, typically measured at a particular Rossland opacity value such as 2⁄3.[6] Photon tired winds appear between both.[6]

A thin proto-galactic disk also surrounds the quasi-star, in which the star feeds it at rates from 2×10−3 to a few of tens of M☉ per year.[8][6] The black hole within the core would accrete rapidly the surrounding gas from the envelope through a convectively or advectively dominated thick accretion disc or via quasi-spherical accretion,[6] with at least 10−4 to several M☉ per year,[8][1] possibly up to a highly super-Eddington accretion rate.[2] The accretion rate onto the black hole energetically sustains the gaseous envelope,[9] but its limit is often set by the Eddington limit of the entire object, which is initially much larger than that of the black hole alone. The excess energy is carried away by convection.[8] Analogous to active galactic nuclei and gamma-ray bursts, the rotation of the quasi-star alongside the poloidal magnetic fields in the black hole or disk magnetosphere (or both) may mediate the production of a relativistic jet, and they may be transported from the outer regions to the center only thanks to thick accretion flows.[1] This would produce gamma rays in the reconfinement shocks formed within one-hundredth to one times the quasi-star's radius.[1]

Properties

edit

Depending on models, quasi-stars by the time of their formation would be at least 1,000 M☉,[10] but they are most likely to be several to dozens millions of M☉, although lower mass quasi-stars may have existed. The most massive possible quasi-stars may reach up to an upper limit for the envelope of about 100 million M☉ based on the assumption that a small degree of rotation can stabilize stars of this mass.[11][6]

They are predicted to have had surface temperatures over 14,000 K (13,700 °C) by the time of their formation,[10][8] with minimum temperatures being of between 5,000 and 3,000 K depending on the metallicity.[10][1] Those estimates usually assume a spherical photosphere; temperatures may be lower due to the star being strongly flattened by its rotation.[2] The envelope and surface of a quasi-star would radiate thermally at the Eddington luminosity and have hence extremely high bolometric luminosities, well over 1011 L☉,[1][8][10] about as luminous as a small galaxy.[4] With such extreme properties, they tend to have radii dwarfing significantly the largest modern stars within the local universe, including red supergiant and asymptotic giant branch stars such as Mu Cephei, VY Canis Majoris, VX Sagittarii, and NML Cygni. By the time of its formation, a quasi-star begins with an initial radius of at least 3,000 R☉. More massive quasi-stars would also result in larger initial radii and luminosities, and as they cool over time, their sizes would increase. Thus, they tend to have maximum radii of at least 22,000 R☉ (100 au),[10] with the largest possible quasi-stars reaching up to at least around 1.4 million R☉ or 0.11 light-years (1×10^15 m) assuming 10 million M☉.[1] Quasi-stars also have extreme mass loss rates through luminosities and winds from their envelopes, in analogy to very massive stars such as Eta Carinae, which are hence surrounded by circumstellar nebulae.[6] They are expected to lose mass up to few to tens of thousands M☉ per year.[6] Thanks to their formation from rotationally supported gas (BVR), most quasi-stars are expected to rotate rapidly by the time of their formation.[10] Furthermore, such rapid rotation would result in the flattening similar to stars like Vega and Achernar, reinforcing their stabilization against dynamical instability such as direct collapse for objects exceeding 100,000 M☉.[10] Over time, they would rotate more slowly due to accretion from the surrounding pre-galactic disc,[12] but still likely faster on the equatorial plane than on the poles, as long as still embedded.[11]

The black hole core inside the star at the time of its formation would have an initial stellar mass between 5 and 100 M☉,[2] but it likely comprises only a small fraction of the quasi-star's core due to rotation.[13] Traditionally, it was expected to have grown to a maximum mass of no more than 1.7 percent of the total mass of a high-mass quasi-star, somewhere at an intermediate or a supermassive mass between about 1,000 M☉ up and several 105 M☉.[10] A more recent research considered those previous estimates for the inner boundary condition for the Bondi radius of the black hole to be an artifact of an unphysical and underestimated formulation, and by using a new model, this study determined that the black hole core surrounded by an inner saturated-convection region could grow to up to 61.8% of the quasi-star's mass, leaving only 36.8% for the saturated-convection region only 2.3% for the outer polytropic envelope.[9]

  • [6]
  • This implies that quasi-stars are unlikely to be significant sources of hard UV radiation when it has grown more than 10,000 M☉. Prior this, the rate of production of ionizing photons would have become very high at 1055 photons, although the total output falls far short of the requirement for reionizing the Universe, given this would last for less than 100,000 years, along with keeping the molecular hydrogen in its surroundings photodissociated. Those estimates correspond to a spherical photosphere, although temperatures of the photosphere could be even lower, under the assumption that the latter is strongly flattened by the star's rotation.[2]

Formation

edit

The formation of quasi-stars could only happen early in the development of the Universe; thus, they may have been very massive Population III and I stars, although earlier research stated they may have existed even before the formation of the first "normal" stars.[2] They may have formed via monolithic collapse of atomic-cooling dark matter halos with a viral temperature over 10,000 K[11][2] or a dense cloud during a protogalaxy collision.[14][6] Drawing in enormous amounts of gas via gravity, this can produce supermassive stars with over tens of thousands of M☉.[15][16][17]

However, the formation of a quasi-star depends on whether the infall of gas is high enough to prevent a thermal equilibrium from being established in the central hydrostatic core.[14] Otherwise, the supermassive star would instead collapse to become a direct collapse black hole.[14] Furthermore, the accretion rate onto the central object may drop over time, and due to the limited gas reservoir in typical dark matter halos (with mass around 107 M☉), the fate of supermassive stars would be more likely to end up as direct collapse black holes. However, a quasi-star may form if high accretion rates of at least 0.14 M☉/yr are maintained within 1–2 Myr.[17][11] This may be possible for a pregalactic halo of at least 6 billion M☉ (for reference, some small galaxies have only 5 million M☉), enough to form a black hole of 104 M☉,[6] or during a merger of puffy protogalaxies or haloes rich in gas and even metal,[8] which could increase the accretion rate up to roughly 105 M☉/yr. This is also only possible if the accreting protostar has already become sufficiently supermassive with a small core; otherwise, it would collapse directly into a dark collapse black hole (a direct collapse without prior stellar phase).[18][19] In addition, the formation of a quasi-star might be disrupted by the presence of a massive black hole already present in the halo.[13]

Although nuclear reactions may begin, the high infall rate proceeds to compress and heat the core until attaining 500 million K, resulting in a catastrophic birth of a black hole seed due to neutrino losses.[10] During the collapse of the stellar progenitor, its inner core forms a stellar-mass black hole and a luminous central accretion disk, initially enclosing a portion of its core inside the envelope supported by radiation pressure.[2][11] Dotan et al.:[6] It would continue generating a large amount of radiant energy from the infall of stellar material. This constant outburst of energy would counteract the force of gravity, creating an equilibrium similar to the one that supports modern fusion-based stars and causing the outer layers of the quasi-star to expand and cool down.[10] X-rays from the BH support the rest of the star from prompt collapse and form a stable envelope that would appear to be a cool, red giant star to an external observer. Such stars can grow to ∼106 M⊙ before the BH becomes so massive that a hydrostatic envelope is no longer possible.[20] The interior black hole may then continue accreting from the stellar envelope. At the same time, the accretion rate of the quasi-star will only be limited by the Eddington rate corresponding to the total mass of the configuration.[17]

  • Begelman et al. 2006:[2]
    • Rapid growth may be limited by feedback from the accretion process and/or disruption of the mass supply by star formation or halo mergers

Models from a 2023 paper, although this study did not follow the collapse of supermassive stars to late times, predicted that it is unlikely that X-rays from the central black hole could halt the collapse of modelled stars (with final masses of (3.5–370)×103 M☉) (need to talk about accretion rates) because of large infall velocities that enclose most of the stellar mass so it will go into the black hole soon after birth. This would thus prevent the formation of a quasi-star that could create black hole seeds of up to a million M☉, and instead become straight direct collapse black holes born with the mass at which their progenitors die.[20]

Evolution

edit
  • [6]
  • [13]
  • QS at the time of its formation
  • Lifespan: Despite those extreme properties, the stabilisation from the central black hole would allow quasi-stars to have a lifespan longer than some very massive normal stars, over 4 million years (Myr), depending on models.[8][10] However, a 2024 model, featuring a saturated-convection region around the black hole core, suggested an additional time of 23 to 46 Myr, during which the black hole would have grown by sixty percent of the quasi-star mass.[9] Low-mass quasi-stars with a central accretion disk inside may have shorter lifespans below 10,000 years before the accretion luminosity dissolves the envelope inside.[11]
  • Super-Eddington: A super-Eddington quasi-star is only possible if all atmospheres become unstable before reaching the Eddington limit; it would last below 1,000 years.[6]
  • The fate of a QS:
    • The fate of a quasi-star remains uncertain. As suggested by many models, as a quasi-star cools over time, its outer envelope would become less opaque once below 10,000 K, until further cooling to a limiting temperature of roughly 5,000–4,000 K (4,730–3,730 °C) for Population III opacities or lower if metal-enriched,[10] as low as roughly 3,000 K for a solar metallicity (e.g Population I stars).[1]
    • Traditionally: Traditionally, earlier models considered the limiting temperature marking the end of the quasi-star's life since there is no hydrostatic equilibrium at or below this limiting temperature.[10] In that case, the object would then quickly dissipate by radiation pressure, leaving behind the central black hole.[10]
  • More things about the photosphere[2]
    • Unlike gas pressure-supported red giants, which remain stable as they hover close to the minimum temperature, QSs lose dynamical equilibrium and disperse once they reach this limit. This is because QSs are supported against collapse by radiation pressure, which also governs internal energy transport – implying that there are not enough degrees of freedom available to allow the envelope to adjust quasi-statically. From equation (4), we see that the BH mass cannot exceed about 1 per cent of the QS mass. For formula⁠, this leads to a final QS mass that substantially exceeds the mass of its supermassive star precursor. For smaller infall rates, the final QS mass will be similar to the mass of the supermassive star (equation 2) and the BH mass will be ≲104 M⊙.
    • At lower temperatures, the Planck mean opacity (which is relevant for calculating the radiation force in LTE, and therefore the Eddington limit) becomes very sensitive to temperature (Mayer & Duschl 2005), increasing sharply at T∗ . 104 K and then decreasing to several orders of magnitude below the electron scattering opacity as the temperature declines further. The sharp decrease in opacity would affect the photosphere at an even earlier stage in the quasistar’s evolution.
    • Similarly, the quasistar produces ≃ 1050 photons s−1 in the Lyman-Werner band, but can keep the molecular hydrogen in its surroundings photodissociated only for 105yr. These estimates correspond to a spherical photosphere at r∗, but we note that photospheric temperatures could be even lower if the photosphere is strongly flattened by rotation. The above estimates are valid only as long as T∗ & 104 K, corresponding to.

Ball

edit

A 2011 study modelled a quasi-star with an initial mass 10,000 M☉ begins its life with an effective temperature of 14,300 K, a luminosity of 3.48×108 L☉ and radius of 3,030 R☉ with a 5 M☉ central black hole accreting at 10−4 M☉/yr.[8] The second feature, apparent in all but the first density profile in Fig. 1, is the density inversion in the outer layers. It appears once the photospheric temperature Tsurf drops below about 8000 K. From then, the surface opacity increases owing to hydrogen recombination. Before the end of the evolution by 3.7 million years, these properties, along with the black hole's accretion rate, achieve their local limits, reaching 4,490 K and 40,700 R☉ with a luminosity of 6.05×108 L☉ and a black hole accretion rate of 3.7×10−4 M☉.[8] By 4.23 million years, the quasi-star's evolution terminates as the black hole reaches a final mass 1,194 M☉ with a cavity mass of 3,360 M☉, but with a decreased accretion rate 3.53×10−4 M☉, with quasi-star's properties being 4,510 K and 39,600 R☉ with a luminosity of 5.81×108 L☉ before hydrostatic equilibrium breaks down. The physical reason for this upper limit remains elusive, but we have made some progress in understanding it using a modified version of the Lane-Emden equation (see Section 4). The existence of the limit is certainly robust as it does not depend on the total mass of the quasi-star over at least two orders of magnitude (see Section 5.4) nor on whether the envelope mass changes in time (see Section 5.3).[8]

Subsequent evolution

edit
  • Whatever happened to the star after is unknown, but it was assumed that the entire object would fall inside the black hole's cavity mass. The material within the cavity, thus the Bondi radius of the black hole, would be already still moving towards the black hole, presumably accreting at its Eddington-limited rate.[8]

Fiacconi

edit
  • [11]
  • In another model, quasi-stars with a relatively low mass might be able to form a central accretion disc and reach an equilibrium configuration but last shorter, for only lesser than few thousands of years before the accretion luminosity unbinds the surrounding envelope, with outflows then suppress the growth of the central black hole. This would result rather a mass for the black hole for 100–1,000 M☉.[11]
    • Within the limitations of this approach (discussed in Section 3.2), we find that, at given M•, most of the massive quasi-stars might not be able to form a central, rotationally-supported accretion region, while the contrary is true for lower mass quasi-stars, typically living within the Evaporation Strip. This bimodal behaviour could lead to different fates, depending on the mass of the original supermassive star at the collapse of the central core that leads to the formation of the central embryo black hole. At high masses, the black hole might swallow most of the mass that is still infalling from larger radii without providing enough feedback either to stabilise the structure or to halt the collapse. The central black hole would then accrete a large fraction of the envelope mass, possibly reaching M• ∼ 104−5 M⊙. On the other hand, less massive envelopes might be able to form a central accretion disc and to reach an equilibrium configuration, i.e. a quasi-star. However, outflows then suppress the growth of the central black hole, leading to M• ∼ 102−3 M⊙.
    • In this paper, we investigate the role of rotation within quasi-stars. Although a fully self-consistent description of the rotating envelope is beyond the purposes of this paper, our treatment of rotation can be used to assess the coupling between the inner accretion region and the massive envelope. Before discussing the implications of our findings for black holes forming inside rotating quasi-stars, we recall their fate, when rotation is not included. This is summarised in Figure 4, adapted from Paper I, to which we refer the reader for more details. Combinations of (M•, M⋆) that lie in the region marked as “No Hydrostatic Solutions” cannot form a stable envelope surrounding the black hole and therefore this latter cannot go through a phase of super-Eddington growth inside a quasi-star. For this to happen, the envelope needs to be at least a few hundred times more massive than its black hole (the parameter space marked as “Growth Region”). There, black holes with an initial mass of ∼ 100 M⊙ can reach in just & 104 yr more than 104 M⊙, depending on the initial envelope mass. This is because the black hole accretes at or beyond the Eddington rate for the envelope mass. Moreover, in these massive envelopes the loss of mass via winds induced by the super-Eddington luminosities proceeds at a lower rate than the black hole growth. The opposite is true for lower envelope masses at the same black hole masses within the “Evaporation Strip”, where outflows remove matter from the envelope faster than black hole accretion, and the latter is then suppressed. Here a quasi-star can form, but it can just last for < 104 yr, with little impact on the embedded black hole. We can now turn our attention to a discussion on how our results might affect this picture.

Direct collapse black hole

edit
  • Saturated-convection and longer lifetime:[9]
    • Here we analyze the intermediate "quasi-star" phase that accompanies some direct collapse models, during which a natal BH accretes mass from and energetically sustains (through accretion) an overlying gaseous envelope. We argue that previous estimates of the maximum BH mass that can be reached during this stage, ∼1% of the total quasi-star mass, are unphysical, and arise from underestimating the efficiency with which energy can be transported outward from regions close to the BH. We construct new quasi-star models that consist of an inner, "saturated-convection" region (which conforms to a convection-dominated accretion flow near the BH) matched to an outer, adiabatic envelope. These solutions exist up to a BH mass of ∼60% the total quasi-star mass, at which point the adiabatic envelope contains only 2\% of the mass (with the remaining ∼38% in the saturated-convection region), and this upper limit is reached within a time of 20−40 Myr. We conclude that quasi-stars remain a viable route for producing SMBHs at large redshifts, consistent with recent JWST observations.
    • , it is significant enough to modify the internal structure of the quasi-star envelope, pushing the hydrostatic region to smaller radii and invalidating the use of the Bondi condition. Rather than implying a limit on BH growth, this bound simply suggests that the energy source responsible for sustaining the overlying envelope is concentrated at smaller radii, closer to the BH. allowing the black hole core to reach over 106 M☉ up to several 107 M☉.
    • Here we analyze the intermediate “quasi-star” phase that accompanies some direct collapse models, during which a natal BH accretes mass from and energetically sustains (through accretion) an overlying gaseous envelope.

Demography

edit
  • [13]
    • The number of quasi-stars forming at redshifts between 2 and 4 has been estimated to be of the order of 0.01 per cubic Mpc in the comoving frame by Volonteri & Begelman (2010) (see their Fig. 5). With this rate we expect a total number of 1010 quasi-stars to form in this redshift range. The number of sources to be observed at a quasi-star stage will be much smaller since this stage lasts thousands to millions of years which is short compared to the age of the universe. A rough estimate of the QS phase duration is given by tQS (see Eq. 5), thus reducing the number of active quasi-stars to a representative value of 3.1 × 105 for all dimensionless parameters equal 1. Since the emission is collimated, the number of sources with jets toward an observer will be further reduced. In general, combining all coefficients we obtain N = 400(θ/4◦)2M −1/8 QS,7 ǫ1/2 d,−2 , (18) which shows that both the typical γ-ray fluxes and the number of the high-latitude, unidentified Fermi sources in the First Fermi LAT Catalog (Abdo et al. 2010) can be reproduced in the QS-jet scenario. Since the number of on-axis QSs at still larger redshifts (from 4 to 10) is only by a factor of 2 larger than the number of resolved sources, their contribution to the still unexplained γ-ray background (Inoue & Totani 2009) can be significant (up to 50%) taking into consideration their bolometric luminosity, but more careful analysis should be done taking into account the source and the background spectral shape. Off-axis sources can contribute to the hard X-ray background, but not significantly (a few per cent or less).
    • Quasi-stars formed in metal-enriched regions may be
    • The most direct prediction of this work is shown in Fig. 4, where we estimate the detectability of QSs in a JWST field of view. At low metallicities, QSs at redshifts of a few may resemble featureless blackbodies, with colours reminiscent of brown dwarfs. Direct redshift measurements may not be feasible, in which case QSs would have to be identified via their massive hosts. Their low numbers – a consequence of their short lifetimes – will make it even more challenging to find them with a telescope having a relatively small field of view. It is possible that some QSs could have formed in metal-enriched regions and even at relatively low redshifts (z≳ 2). These might be relatively easy to detect but extremely rare. Other, secondary characteristics of QSs might conceivably aid in their detection, e.g. if accretion on to the BH deep in the core leads to the formation of a jet that pierces the stellar surface.
  • [1]
  • [2]
    • Both analytic and numerical models, computed using Pop III opacities, suggest that this minimum temperature is around 4000- 5000 K. If quasistars are implicated in the formation of seeds for supermassive black holes in pre-galactic haloes, as suggested by Begelman, Volonteri & Rees (2006), this floor temperature implies that the most luminous quasistars would emit most strongly in the rest-frame near-IR. At typical redshifts of z ≃ 10, the observed spectrum would peak at λ ∼ 10μm.

Possible detection

edit
  • As such, researches also proposed that quasi-star jets may be accounted for a large fraction of unidentified gamma-ray sources located at high latitudes, in which most of them are considered to be extragalactic.[1] Like blazars, they would produce nonthermal spectra at lower energies (optical-IR) dominated by the synchrotron mechanism and in the gamma-ray band by the inverse-Compton process. However, they can be distinguished from blazars depending on the ratio of the gamma-ray to the IR components and the presence of broad emission lines. Most of BL Lacertae objects are expected to have a low ratio and no broad emission lines.[1]
  • The most direct prediction of this work is shown in Fig. 4, where we estimate the detectability of QSs in a JWST field of view. At low metallicities, QSs at redshifts of a few may resemble featureless blackbodies, with colours reminiscent of brown dwarfs. Direct redshift measurements may not be feasible, in which case QSs would have to be identified via their massive hosts. Their low numbers – a consequence of their short lifetimes – will make it even more challenging to find them with a telescope having a relatively small field of view. It is possible that some QSs could have formed in metal-enriched regions and even at relatively low redshifts (z≳ 2). These might be relatively easy to detect but extremely rare. Other, secondary characteristics of QSs might conceivably aid in their detection, e.g. if accretion on to the BH deep in the core leads to the formation of a jet that pierces the stellar surface.[13]
  • Our results can be summarized as follows.[13]
    • (i)A limited range of supermassive star/QS/MBH formation efficiencies exists that allows one to reproduce observational constraints (the density of z= 6 quasars, the cumulative mass density accreted on to MBHs, the current mass density of MBHs and reionization). These constraints translate into 0.01 ≤λthr≤ 0.02.
    • (ii)The mass function of QSs peaks at MQS≃ 106 M⊙ and decreases almost as a power law with slope ≃−(1.5–2). The more efficient QS and MBH formation is, the steeper the mass function. This is because more seeds form at the highest redshifts when the hosts have relatively low mass and circular velocity.
    • (iii)Modelling QS emission as a blackbody at TQS= 4000 K, we calculate the number counts for the JWST 2–10 μm band. Assuming a sensitivity of 10 nJy at 2 μm, we find that JWST could detect up to several QSs per field at z≃ 5–10.
    • (iv)The redshift of formation of QSs increases with the cosmic bias of their hosts. We therefore expect that the latest forming QSs should be found in the field rather than in a high-density environment.
    • (v)The number density of MBHs can be dominated by either the descendants of Pop III remnants or QS MBHs. In the low-efficiency case (λthr= 0.01) the number density of MBHs is dominated by the Pop III channel, if available; the opposite is true for the high-efficiency case (λthr= 0.02). In contrast to the number density of MBHs, the mass density is always dominated by QS MBHs.
    • (vi)If the only channel of MBH formation is via QS seeds, then the mass function of MBHs cuts off below log(MBH/ M⊙) = 4.5, where the seed mass function drops as well. If Pop III remnants offer an alternative route to MBH formation, then the mass function is double-peaked, with each peak tracing a different seed formation mechanism. The more efficient QS formation is, the more pronounced the peak at log(MBH/M⊙) ≃ 4.5 becomes.[13]

See also

edit

References

edit
  1. ^ a b c d e f g h i j k l m Czerny, Bozena; Janiuk, Agnieszka; Sikora, Marek; Lasota, Jean-Pierre (2012). "Quasi-Star Jets as Unidentified Gamma-Ray Sources". The Astrophysical Journal. 755 (1): L15. arXiv:1207.1560. Bibcode:2012ApJ...755L..15C. doi:10.1088/2041-8205/755/1/L15. S2CID 113397287.
  2. ^ a b c d e f g h i j k l m Begelman, M. C.; et al. (Jun 2006). "Formation of supermassive black holes by direct collapse in pre-galactic haloed". Monthly Notices of the Royal Astronomical Society. 370 (1): 289–298. arXiv:astro-ph/0602363. Bibcode:2006MNRAS.370..289B. doi:10.1111/j.1365-2966.2006.10467.x. S2CID 14545390.
  3. ^ . arXiv:2507.09085. {{cite arXiv}}: Missing or empty |title= (help)CS1 maint: missing class (link) A bot will complete this citation soon. Click here to jump the queue
  4. ^ a b c Battersby, Stephen (29 November 2007). "Biggest black holes may grow inside 'quasistars'". New Scientist.
  5. ^ Clery, Daniel (29 July 2025). "Early universe's 'little red dots' may be black hole stars".
  6. ^ a b c d e f g h i j k l m n o p q Dotan, Calanit; Rossi, Elena M.; Shaviv, Nir J. (2011). "A lower limit on the halo mass to form supermassive black holes". Monthly Notices of the Royal Astronomical Society. 417 (4): 3035–3046. arXiv:1107.3562. Bibcode:2011MNRAS.417.3035D. doi:10.1111/j.1365-2966.2011.19461.x.
  7. ^ Cite error: The named reference vonlonteri was invoked but never defined (see the help page).
  8. ^ a b c d e f g h i j k l Ball, Warrick H.; Tout, Christopher A.; Żytkow, Anna N.; Eldridge, John J. (2011-07-01). "The structure and evolution of quasi-stars: The structure and evolution of quasi-stars". Monthly Notices of the Royal Astronomical Society. 414 (3): 2751–2762. arXiv:1102.5098. Bibcode:2011MNRAS.414.2751B. doi:10.1111/j.1365-2966.2011.18591.x.
  9. ^ a b c d e f . arXiv:2405.00084. {{cite arXiv}}: Missing or empty |title= (help)CS1 maint: missing class (link) A bot will complete this citation soon. Click here to jump the queue
  10. ^ a b c d e f g h i j k l m n Begelman, Mitch; Rossi, Elena; Armitage, Philip (2008). "Quasi-stars: accreting black holes inside massive envelopes". MNRAS. 387 (4): 1649–1659. arXiv:0711.4078. Bibcode:2008MNRAS.387.1649B. doi:10.1111/j.1365-2966.2008.13344.x. S2CID 12044015.
  11. ^ a b c d e f g h Fiacconi, Davide; Rossi, Elena M. (2017). "Light or heavy supermassive black hole seeds: The role of internal rotation in the fate of supermassive stars". Monthly Notices of the Royal Astronomical Society. 464 (2): 2259–2269. arXiv:1604.03936. doi:10.1093/mnras/stw2505.{{cite journal}}: CS1 maint: unflagged free DOI (link)
  12. ^ Cite error: The named reference dolan was invoked but never defined (see the help page).
  13. ^ a b c d e f g . doi:10.1111/j.1365-2966.2010.17359.x. {{cite journal}}: Cite journal requires |journal= (help); Missing or empty |title= (help)
  14. ^ a b c Ball, Warrick H. (2012). "Quasi-stars and the Schönberg-Chandrasekhar limit". arXiv:1207.5972.{{cite arXiv}}: CS1 maint: missing class (link) A bot will complete this citation soon. Click here to jump the queue
  15. ^ Yasemin Saplakoglu (September 29, 2017). "Zeroing In on How Supermassive Black Holes Formed". Scientific American. Retrieved April 8, 2019.
  16. ^ Mara Johnson-Goh (November 20, 2017). "Cooking up supermassive black holes in the early universe". Astronomy. Retrieved April 8, 2019.
  17. ^ a b c Schleicher, Dominik R. G.; Palla, Francesco; Ferrara, Andrea; Galli, Daniele; Latif, Muhammad (25 May 2013). "Massive black hole factories: Supermassive and quasi-star formation in primordial halos". Astronomy & Astrophysics. 558: A59. arXiv:1305.5923. Bibcode:2013A&A...558A..59S. doi:10.1051/0004-6361/201321949. S2CID 119197147.
  18. ^ Mayer, Lucio; Bonoli, Silvia (2019). "The route to massive black hole formation via merger-driven direct collapse: A review". Reports on Progress in Physics. 82 (1): 016901. arXiv:1803.06391. Bibcode:2019RPPh...82a6901M. doi:10.1088/1361-6633/aad6a5. PMID 30057369. S2CID 51865966.
  19. ^ Haemmerlé, L.; Meynet, G.; Mayer, L.; Klessen, R. S.; Woods, T. E.; Heger, A. (2019). "Maximally accreting supermassive stars: A fundamental limit imposed by hydrostatic equilibrium". Astronomy & Astrophysics. 632: L2. arXiv:1910.04776. Bibcode:2019A&A...632L...2H. doi:10.1051/0004-6361/201936716.
  20. ^ a b Herrington, Nicholas P.; Whalen, Daniel J.; Woods, Tyrone E. (2023). "Modelling supermassive primordial stars with <SCP>mesa</SCP>". Monthly Notices of the Royal Astronomical Society. 521: 463–473. doi:10.1093/mnras/stad572.
edit
Retrieved from "https://en.wikipedia.org/w/index.php?title=User:ZaperaWiki44/sandbox/Quasi-star&oldid=1308650312"