Tolman–Oppenheimer–Volkoff limit

The Tolman–Oppenheimer–Volkoff limit (or TOV limit) is an upper bound to the mass of cold, non-rotating neutron stars, analogous to the Chandrasekhar limit for white dwarf stars. Stars more massive than the TOV limit collapse into a black hole. The original calculation in 1939, which neglected complications such as nuclear forces between neutrons, placed this limit at approximately 0.7 solar masses (M_☉). Later, more refined analyses have resulted in larger values.

Theoretical work in 1996 placed the limit at approximately 1.5 to 3.0 M_☉,^[1] corresponding to an original stellar mass of 15 to 20 M_☉; additional work in the same year gave a more precise range of 2.2 to 2.9 M_☉.^[2]

Data from GW170817, the first gravitational wave observation attributed to merging neutron stars (thought to have collapsed into a black hole^[3] within a few seconds after merging^[4]) placed the limit in the range of 2.01 to 2.17 M_☉.^[5]

In the case of a rigidly spinning neutron star, meaning that different levels in the interior of the star all rotate at the same rate, the mass limit is thought to increase by up to 18–20%.^[4][5]

History

The idea that there should be an absolute upper limit for the mass of a cold (as distinct from thermal pressure supported) self-gravitating body dates back to the 1932 work of

given roughly by

${\displaystyle P={\frac {1}{\lambda ^{4}}}}$.

At the upper mass limit, that pressure will equal the pressure needed to resist gravity. The pressure to resist gravity for a body of mass M will be given according to the virial theorem roughly by

where ρ is the density. This will be given by ρ = m/λ³, where m is the relevant mass per particle. It can be seen that the wavelength cancels out so that one obtains an approximate mass limit formula of the very simple form

In this relationship, m can be taken to be given roughly by the

) for which the fermionic particles providing the pressure are electrons. This is because the mass density is provided by the nuclei in which the neutrons are at most about as numerous as the protons. Likewise the protons, for charge neutrality, must be exactly as numerous as the electrons outside.

In the case of

for white dwarfs. Oppenheimer and Volkoff's paper noted 'If... the effect of repulsive forces, i.e., of raising the pressure for a given density above the value given by the Fermi equation of state, this could tend to prevent the collapse'.[7] And indeed, the most massive neutron star detected so far, PSR J0952–0607, is estimated to be much heavier than Oppenheimer and Volkoff's TOV limit at 2.35±0.17 M_☉.^[8][9] More realistic models neutron stars including baryon strong force repulsion predict a neutron star mass limit of 2.2 to 2.9 M_☉.^[10][11] The uncertainty in the value reflects the fact that the equations of state for extremely dense matter are not well known.

Applications

In a star less massive than the limit, the gravitational compression is balanced by short-range repulsive neutron–neutron interactions mediated by the strong force and also by the quantum degeneracy pressure of neutrons, preventing collapse.^[12]: 74 If its mass is above the limit, the star will collapse to some denser form. It could form a black hole, or change composition and be supported in some other way (for example, by quark degeneracy pressure if it becomes a quark star). Because the properties of hypothetical, more exotic forms of degenerate matter are even more poorly known than those of neutron-degenerate matter, most astrophysicists assume, in the absence of evidence to the contrary, that a neutron star above the limit collapses directly into a black hole.

A

black hole mergers involving black holes in the 7.5–50 solar mass range; it is possible – although unlikely – that these black holes were themselves the result of previous mergers.

Oppenheimer and Volkoff discounted the influence of heat, stating in reference to work by Landau (1932), 'even [at] 10⁷ degrees... the pressure is determined essentially by the density only and not by the temperature'[7] - yet it has been estimated^[16] that temperatures can reach up to approximately >10⁹ K during formation of a neutron star, mergers and binary accretion. Another source of heat and therefore collapse-resisting pressure in neutron stars is 'viscous friction in the presence of differential rotation.'^[16]

Oppenheimer and Volkoff's calculation of the mass limit of neutron stars also neglected to consider the rotation of neutron stars, however we now know that neutron stars are capable of spinning at much faster rates than were known in Oppenheimer and Volkoff's time. The fastest-spinning neutron star known is PSR J1748-2446ad, rotating at a rate of 716 times per second^[17][18] or 43,000 revolutions per minute, giving a linear (tangential) speed at the surface on the order of 0.24c (i.e., nearly a quarter the speed of light). Star rotation interferes with convective heat loss during supernova collapse, so rotating stars are more likely to collapse directly to form a black hole ^[19]: 1044

List of the most massive neutron stars

Below is a list of neutron stars. These include rotating neutron stars and thus are not directly related to the TOV Limit.

Name	Mass (M☉)	Distance ( ly )	Companion class	Mass determination method	Notes	Refs.
PSR J1748-2021B	2.548+0.047 −0.078	27,700	D	Rate of advance of periastron.	In globular cluster NGC 6440.	[20][21][22][23]
4U 1700-37	2.44±0.27	6,910 ± 1,120	O6.5Iaf+	Monte Carlo simulations of thermal comptonization process.	HMXB system.	[24][25]
PSR J0952–0607	2.35±0.17	3,200-5,700			Fastest and heaviest known galactic neutron star	[26]
PSR J1311–3430	2.15–2.7	6,500–12,700	Substellar object	Spectroscopic and photometric observation.	Black widow pulsar.	[27][28]
PSR J1600−3053	2.3+0.7 −0.6	6,500 ± 1,000	D	Fourier analysis of Shapiro delay's orthometric ratio.		[29][30]
PSR J2215+5135	2.27+0.17 −0.15	10,000	G5V	Innovative measurement of companion's radial velocity.	Redback pulsar.	[31]
XMMU J013236.7+303228	2.2+0.8 −0.6	2,730,000	B1.5IV	Detailed spectroscopic modelling.	In M33, HMXB system.	[32]
PSR J0751+1807	2.10±0.2	6,500 ± 1,300	D	Precision pulse timing measurements of relativistic orbital decay.		[33]
PSR J0740+6620	2.08±0.07	4,600	D	Range and shape parameter of Shapiro delay.	Most massive neutron star with a well-constrained mass	[34][35][36]
PSR J0348+0432	2.01±0.04	2,100	D	Spectroscopic observation and orbital decay due to radiation of gravitational waves.		[29][37]
PSR B1516+02B	1.94+0.17 −0.19	24,500	D	Rate of advance of periastron.	In globular cluster M5.	[29][38]
PSR J1614−2230	1.908±0.016	3,900	D	Range and shape parameter of Shapiro delay.	In Milky Way's galactic disk .	[29][30][39]
Vela X-1	1.88±0.13	6,200 ± 650	B0.5Ib	Rate of advance of periastron.	Prototypical detached HMXB system.	[40]
PSR B1957+20	1.81±0.17	6,500	Substellar object	Rate of advance of periastron.	Prototype star of black widow pulsars.	[41][42][43]

List of least massive black holes

Below is a list of black holes.

Name	Mass (M☉)	Distance ( ly )	Companion class	Mass determination method	Notes	Refs.
V723 Monocerotis	3.04±0.06	1,500	K0/K1III	Spectroscopic radial velocity measurements of companion.	Mass may be underestimated due not accurate measurement distance to the companion star.[clarify]	[44]
2MASS J05215658+4359220	3.3+2.8 −0.7	10,000	K-type (?) giant	Spectroscopic radial velocity measurements of noninteracting companion.	In Milky Way outskirts.	[29][45][46]
GW190425's remnant	3.4+0.3 −0.1	518,600,000	N/A	Gravitational wave data of neutron star merger from LIGO and Virgo interferometers.	97% chance of prompt collapse into a black hole immediately after merger.	[29][47][48]
NGC 3201-1	4.36±0.41	15,600	(see Notes)	Spectroscopic radial velocity measurements of noninteracting companion.	In globular cluster main sequence turn-off .	[29][49]
HR 6819 (QV Tel)	5.0±0.4	1,120	Be/B3III	Spectroscopic radial velocity measurements of companion.	Unconfirmed black hole.	[50]
GRO J1719-24/ GRS 1716−249	≥4.9	8,500	K0-5 V	Near-infrared photometry of companion and Eddington flux.	LMXB system.	[29][51]
4U 1543-47	5.0+2.5 −2.3	30,000 ± 3,500	A2 (V?)	Spectroscopic radial velocity measurements of companion.	SXT system.	[29][52]
XTE J1650-500	≥5.1	8,500 ± 2,300	K4V	QPOs	Transient binary X-ray source	[53][54]
GRO J1655-40	5.31±0.07	<5,500	F6IV	Precision X-ray timing observations from RossiXTE.	LMXB system.	[55][56]
GX 339-4	5.9±3.6	26,000	N/A			[29]

List of objects in mass gap

This list contains objects that may be neutron stars, black holes, quark stars, or other exotic objects. This list is distinct from the list of least massive black holes due to the undetermined nature of these objects, largely because of indeterminate mass, or other poor observation data.

Name	Mass (M☉)	Distance ( ly )	Companion class	Mass determination method	Notes	Refs.
GW170817's remnant	2.74+0.04 −0.01	144,000,000	N/A	Gravitational wave data of neutron star merger from LIGO and Virgo interferometers.	In NGC 4993. Possibly collapsed into a black hole 5–10 seconds after merger.	[57]
SS 433	3.0–30.0	18,000 ± 700	A7Ib		First discovered microquasar system. Confirmed to have a magnetic field, which is atypical for a black hole; however, it could be the field of the accretion disk, not of the compact object.	[58][59][60]
LB-1	2.0–70.0	approx. 7,000	Be star/stripped helium star		Initially thought to be first black hole in pair-instability mass gap.	[61][62]
Cygnus X-3	2.0–5.0	24,100 ± 3,600	WN4-6	Near-infrared spectroscopy and atmosphere model fitting of companion.	Microquasar system. Major differences between the spectrum of Cyg X-3 and typical accreting BH can be explained by properties of its companion star.	[63][64]
LS I +61 303	1.0 - 4.0	7,000	B0Ve	Spectroscopic radial velocity measurements of companion.	Microquasar system. It has a spectrum typical for black holes, however it emits HE and VHE gamma rays similar to neutron stars LS_2883 and HESS J0632+057, as well as mysterious object LS 5039.	[65][66]
LS 5039	3.7+1.3 −1.0	8,200 ± 300	O(f)N6.5V	Intermediate-dispersion spectroscopy and atmosphere model fitting of companion.	Microquasar system. Only lowest possible mass allows it not to be a black hole.	[67]
V518 Persei	3.97+0.95 −1.87	8,500	M4.5V	Photometric light curve modelling.	SXT system. Only mass close to lowest possible allows it not to be a black hole.	[29][68]

See also

• Tolman–Oppenheimer–Volkoff equation
• Bekenstein bound
• Quark star

Notes

References