3.5 Supermassive stars

There is a large body of observational evidence that supermassive black holes (SMBHs, M ≳ 106M ⊙) exist in the centers of many, if not most galaxies (see, e.g., the reviews of Rees [246] and Macchetto [188]). The masses of SMBHs in the centers of more than 45 galaxies have been estimated from observations [83] and there are more than 30 galaxies in which the presence of a SMBH has been confirmed [163].

3.5.1 Supermassive stars: Evolution

One of the possible formation mechanisms for SMBHs involves the gravitational collapse of SMSs. The timescale for this formation channel is short enough to account for the presence of SMBHs at redshifts z > 6 [155Jump To The Next Citation Point]. SMSs may contract directly out of the primordial gas, if radiation and/or magnetic-field pressure prevent fragmentation [131, 75, 130, 186Jump To The Next Citation Point, 29, 1, 348Jump To The Next Citation Point]. Alternatively, they may build up from fragments of stellar collisions in clusters [262, 14]. Supermassive stars are radiation dominated, isentropic and convective [276Jump To The Next Citation Point, 345, 186]. Thus, they are well represented by an n = 3 polytrope. If the star’s mass exceeds 6 10 M ⊙, nuclear burning and electron/positron annihilation are not important.

After formation, an SMS will evolve through a phase of quasi-stationary cooling and contraction. If the SMS is rotating when it forms, conservation of angular momentum requires that it spins up as it contracts. There are two possible evolutionary regimes for a cooling SMS. The path taken by an SMS depends on the strength of its viscosity and magnetic fields and on the nature of its angular momentum distribution.

In the first regime, viscosity or magnetic fields are strong enough to enforce uniform rotation throughout the star as it contracts. Baumgarte and Shapiro [12Jump To The Next Citation Point] have studied the evolution of a uniformly-rotating SMS up to the onset of relativistic instability. They demonstrated that a uniformly-rotating, cooling SMS will eventually spin up to its mass shedding limit. The mass shedding limit is encountered when matter at the star’s equator rotates with the Keplerian velocity. The limit can be represented as βshed = (T∕|W |)shed. In this case, βshed = 9 × 10−3. The star will then evolve along a mass-shedding sequence, losing both mass and angular momentum. It will eventually contract to the onset of relativistic instability [151, 49, 50, 276Jump To The Next Citation Point, 155, 185Jump To The Next Citation Point].

The SMS formation scenario for SMBHs is just one of many. It fits into a broad class of scenarios invoking the collapse of supermassive objects formed in halos of dense gas, e.g., [15]. Although the structure used in many of these SMS calculations may not be appropriate for this broad class of supermassive objects, many of the basic features studied in these SMS simulations will persist.

Baumgarte and Shapiro used both a second-order, post-Newtonian approximation and a fully general-relativistic numerical code to determine that the onset of relativistic instability occurs at a ratio of R ∕M ∼ 450, where R is the star’s radius and G = c = 1 in the remainder of this section. Note that a second-order, post-Newtonian approximation was needed because rotation stabilizes the destabilizing role of non-linear gravity at the first post-Newtonian level. If the mass of the star exceeds 106M ⊙, the star will then collapse and possibly form a SMBH. If the star is less massive, nuclear reactions may lead to explosion instead of collapse.

The major result of Baumgarte and Shapiro’s work is that the universal values of the following ratios exist for the critical configuration at the onset of relativistic instability: T∕|W |, R∕M, and J∕M 2. These ratios are completely independent of the mass of the star or its prior evolution. Because uniformly-rotating SMSs will begin to collapse from a universal configuration, the subsequent collapse and the resulting gravitational waveform will be unique.

In the opposite evolutionary regime, neither viscosity nor magnetic fields are strong enough to enforce uniform rotation throughout the cooling SMS as it contracts. In this case, it has been shown that the angular momentum distribution is conserved on cylinders during contraction [22]. Because viscosity and magnetic fields are weak, there is no means of redistributing angular momentum in the star. So, even if the star starts out rotating uniformly, it cannot remain so.

The star will then rotate differentially as it cools and contracts. In this case, the subsequent evolution depends on the star’s initial angular momentum distribution, which is largely unknown. One possible outcome is that the star will spin up to mass-shedding (at a different value of βshed than a uniformly-rotating star) and then follow an evolutionary path that may be similar to that described by Baumgarte and Shapiro [12Jump To The Next Citation Point]. The alternative outcome is that the star will encounter the dynamical bar instability prior to reaching the mass-shedding limit. New and Shapiro [222, 223] have demonstrated that a bar-mode phase is likely to be encountered by differentially-rotating SMSs with a wide range of initial angular-momentum distributions. This mode will transport mass and angular momentum outward and thus may hasten the onset of collapse.

3.5.2 Rates of supermassive stars

An estimate of the rate of the collapse of SMSs can be derived from the quasar luminosity function. Haehnelt [129] has used the quasar luminosity function to compute the rate of GW bursts from SMBHs, assuming that each quasar emits one such burst during its lifetime (and that each quasar is a supermassive black hole). If it is assumed that each of these bursts is due to the formation of a supermassive black hole via the collapse of an SMS, then Haehnelt’s rate estimates can be used as estimates of the rate of SMS collapse. This rate is likely an overestimate of the SMS collapse rate because many SMBHs may have been formed via merger. Haehnelt predicts that the integrated event rate through redshift z = 4.5 ranges from ∼ 10−6 yr−1 for M = 108M ⊙ objects to ∼ 1 yr−1 for M = 106M ⊙ objects. Thus, as in the case of Population III stars, a reasonable occurrence rate can be determined for an observation (luminosity) distance of 50 Gpc.

3.5.3 Gravitational-wave emission mechanisms of supermassive stars

The GW emission mechanisms related to the collapse of SMSs are a subset of those discussed in the sections on AIC, SNe/collapsars, and Population III stellar collapse.

Bounce: Like very massive stars [97], it is likely that these SMSs will first form a proto black hole before collapsing to form a black hole. If so, they can emit GWs from the bounce that occurs when the proto black hole is formed. Because of the mass of the SMS, an aspherical collapse can produce a much stronger signal than any normal supernovae.

Post-Bounce: GWs may be emitted due to global rotational and fragmentation instabilities that may arise during the collapse/explosion and in the collapsed remnant (prior to black-hole formation).

Black Hole: Ringing in the black hole as it forms and accretes matter may also drive a strong signal.

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