Neutron star

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https://en.wikipedia.org/wiki/Neutron_star

Central neutron star at the heart of the Crab Nebula

Radiation from the rapidly spinning pulsar PSR B1509-58 makes nearby gas emit X-rays (gold) and illuminates the rest of the nebula, here seen in infrared (blue and red).

A neutron star is the collapsed core of a massive supergiant star. It results from the supernova explosion of a massive star—combined with gravitational collapse—that compresses the core past white dwarf star density to that of atomic nuclei. Surpassed only by black holes, neutron stars are the second smallest and densest known class of stellar objects. Neutron stars have a radius on the order of 10 kilometers (6 mi) and a mass of about 1.4 M_☉. Stars that collapse into neutron stars have a total mass of between 10 and 25 solar masses (M_☉), or possibly more for those that are especially rich in elements heavier than hydrogen and helium.

Once formed, neutron stars no longer actively generate heat and cool over time, but they may still evolve further through collisions or accretion. Most of the basic models for these objects imply that they are composed almost entirely of neutrons, as the extreme pressure causes the electrons and protons present in normal matter to combine into additional neutrons. These stars are partially supported against further collapse by neutron degeneracy pressure, just as white dwarfs are supported against collapse by electron degeneracy pressure. However, this is not by itself sufficient to hold up an object beyond 0.7 M_☉ and repulsive nuclear forces increasingly contribute to supporting more massive neutron stars. If the remnant star has a mass exceeding the Tolman–Oppenheimer–Volkoff limit, which ranges from 2.2–2.9 M_☉, the combination of degeneracy pressure and nuclear forces is insufficient to support the neutron star, causing it to collapse and form a black hole. The most massive neutron star detected so far, PSR J0952–0607, is estimated to be 2.35±0.17 M_☉.

Newly formed neutron stars may have surface temperatures of ten million K or more. However, since neutron stars generate no new heat through fusion, they inexorably cool down after their formation. Consequently, a given neutron star reaches a surface temperature of one million K when it is between one thousand and one million years old. Older and even-cooler neutron stars are still easy to discover. For example, the well-studied neutron star, RX J1856.5−3754, has an average surface temperature of about 434,000 K. For comparison, the Sun has an effective surface temperature of 5,780 K.

Neutron star material is remarkably dense: a normal-sized matchbox containing neutron-star material would have a weight of approximately 3 billion tonnes, the same weight as a 0.5-cubic-kilometer chunk of the Earth (a cube with edges of about 800 meters) from Earth's surface.

As a star's core collapses, its rotation rate increases due to conservation of angular momentum, so newly formed neutron stars typically rotate at up to several hundred times per second. Some neutron stars emit beams of electromagnetic radiation that make them detectable as pulsars, and the discovery of pulsars by Jocelyn Bell Burnell and Antony Hewish in 1967 was the first observational suggestion that neutron stars exist. The fastest-spinning neutron star known is PSR J1748−2446ad, rotating at a rate of 716 times per second 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).

There are thought to be around one billion neutron stars in the Milky Way, and at a minimum several hundred million, a figure obtained by estimating the number of stars that have undergone supernova explosions. However, many of them have existed for a long period of time and have cooled down considerably. These stars radiate very little electromagnetic radiation; most neutron stars that have been detected occur only in certain situations in which they do radiate, such as if they are a pulsar or a part of a binary system. Slow-rotating and non-accreting neutron stars are difficult to detect, due to the absence of electromagnetic radiation; however, since the Hubble Space Telescope's detection of RX J1856.5−3754 in the 1990s, a few nearby neutron stars that appear to emit only thermal radiation have been detected.

Neutron stars in binary systems can undergo accretion, in which case they emit large amounts of X-rays. During this process, matter is deposited on the surface of the stars, forming "hotspots" that can be sporadically identified as X-ray pulsar systems. Additionally, such accretions are able to "recycle" old pulsars, causing them to gain mass and rotate extremely quickly, forming millisecond pulsars. Furthermore, binary systems such as these continue to evolve, with many companions eventually becoming compact objects such as white dwarfs or neutron stars themselves, though other possibilities include a complete destruction of the companion through ablation or collision.

The study of neutron star systems is central to gravitational wave astronomy. The merger of binary neutron stars produces gravitational waves and may be associated with kilonovae and short-duration gamma-ray bursts. In 2017, the LIGO and Virgo interferometer sites observed GW170817, the first direct detection of gravitational waves from such an event. Prior to this, indirect evidence for gravitational waves was inferred by studying the gravity radiated from the orbital decay of a different type of (unmerged) binary neutron system, the Hulse–Taylor pulsar.

Formation

Simplified representation of the formation of neutron stars

Any main-sequence star with an initial mass of greater than 8 M_☉ (eight times the mass of the Sun) has the potential to become a neutron star. As the star evolves away from the main sequence, stellar nucleosynthesis produces an iron-rich core. When all nuclear fuel in the core has been exhausted, the core must be supported by degeneracy pressure alone. Further deposits of mass from shell burning cause the core to exceed the Chandrasekhar limit. Electron-degeneracy pressure is overcome, and the core collapses further, causing temperatures to rise to over 5×10⁹ K (5 billion K). At these temperatures, photodisintegration (the breakdown of iron nuclei into alpha particles due to high-energy gamma rays) occurs. As the temperature of the core continues to rise, electrons and protons combine to form neutrons via electron capture, releasing a flood of neutrinos. When densities reach a nuclear density of 4×10¹⁷ kg/m³, a combination of strong force repulsion and neutron degeneracy pressure halts the contraction. The contracting outer envelope of the star is halted and rapidly flung outwards by a flux of neutrinos produced in the creation of the neutrons, resulting in a supernova and leaving behind a neutron star. However, if the remnant has a mass greater than about 3 M_☉, it instead becomes a black hole.

As the core of a massive star is compressed during a Type II supernova or a Type Ib or Type Ic supernova, and collapses into a neutron star, it retains most of its angular momentum. Because it has only a tiny fraction of its parent's radius (sharply reducing its moment of inertia), a neutron star is formed with very high rotation speed and then, over a very long period, it slows. Neutron stars are known that have rotation periods from about 1.4 ms to 30 s. The neutron star's density also gives it very high surface gravity, with typical values ranging from 10¹² to 10¹³ m/s² (more than 10¹¹ times that of Earth). One measure of such immense gravity is the fact that neutron stars have an escape velocity of over half the speed of light. The neutron star's gravity accelerates infalling matter to tremendous speed, and tidal forces near the surface can cause spaghettification.

Properties

Equation of state

The equation of state of neutron stars is not currently known. This is because neutron stars are the second most dense known object in the universe, only less dense than black holes. The extreme density means there is no way to replicate the material on Earth in laboratories, which is how equations of state for other things like ideal gases are tested. The closest neutron star is many parsecs away, meaning there is no feasible way to study it directly. While it is known neutron stars should be similar to a degenerate gas, it cannot be modeled strictly like one (as white dwarfs are) because of the extreme gravity. General relativity must be considered for the neutron star equation of state because Newtonian gravity is no longer sufficient in those conditions. Effects such as quantum chromodynamics (QCD), superconductivity, and superfluidity must also be considered.

At the extraordinarily high densities of neutron stars, ordinary matter is squeezed to nuclear densities. Specifically, the matter ranges from nuclei embedded in a sea of electrons at low densities in the outer crust, to increasingly neutron-rich structures in the inner crust, to the extremely neutron-rich uniform matter in the outer core, and possibly exotic states of matter at high densities in the inner core.

Understanding the nature of the matter present in the various layers of neutron stars, and the phase transitions that occur at the boundaries of the layers is a major unsolved problem in fundamental physics. The neutron star equation of state encodes information about the structure of a neutron star and thus tells us how matter behaves at the extreme densities found inside neutron stars. Constraints on the neutron star equation of state would then provide constraints on how the strong force of the standard model works, which would have profound implications for nuclear and atomic physics. This makes neutron stars natural laboratories for probing fundamental physics.

For example, the exotic states that may be found at the cores of neutron stars are types of QCD matter. At the extreme densities at the centers of neutron stars, neutrons become disrupted giving rise to a sea of quarks. This matter's equation of state is governed by the laws of quantum chromodynamics and since QCD matter cannot be produced in any laboratory on Earth, most of the current knowledge about it is only theoretical.

Different equations of state lead to different values of observable quantities. While the equation of state is only directly relating the density and pressure, it also leads to calculating observables like the speed of sound, mass, radius, and Love numbers. Because the equation of state is unknown, there are many proposed ones, such as FPS, UU, APR, L, and SLy, and it is an active area of research. Different factors can be considered when creating the equation of state such as phase transitions.

Another aspect of the equation of state is whether it is a soft or stiff equation of state. This relates to how much pressure there is at a certain energy density, and often corresponds to phase transitions. When the material is about to go through a phase transition, the pressure will tend to increase until it shifts into a more comfortable state of matter. A soft equation of state would have a gently rising pressure versus energy density while a stiff one would have a sharper rise in pressure. In neutron stars, nuclear physicists are still testing whether the equation of state should be stiff or soft, and sometimes it changes within individual equations of state depending on the phase transitions within the model. This is referred to as the equation of state stiffening or softening, depending on the previous behavior. Since it is unknown what neutron stars are made of, there is room for different phases of matter to be explored within the equation of state.

Density and pressure

Comparison of a 10 km radius neutron star (top left corner) and a 6000 km radius white dwarf, the latter roughly the size of Earth.

Neutron stars have overall densities of 3.7×10¹⁷ to 5.9×10¹⁷ kg/m³ (2.6×10¹⁴ to 4.1×10¹⁴ times the density of the Sun), which is comparable to the approximate density of an atomic nucleus of 3×10¹⁷ kg/m³. The density increases with depth, varying from about 1×10⁹ kg/m³ at the crust to an estimated 6×10¹⁷ or 8×10¹⁷ kg/m³ deeper inside. Pressure increases accordingly, from about 3.2×10³¹ Pa (32 QPa) at the inner crust to 1.6×10³⁴ Pa in the center.

A neutron star is so dense that one teaspoon (5 milliliters) of its material would have a mass over 5.5×10¹² kg, about 900 times the mass of the Great Pyramid of Giza. The entire mass of the Earth at neutron star density would fit into a sphere 305 m in diameter, about the size of the Arecibo Telescope.

In popular scientific writing, neutron stars are sometimes described as macroscopic atomic nuclei. Indeed, both states are composed of nucleons, and they share a similar density to within an order of magnitude. However, in other respects, neutron stars and atomic nuclei are quite different. A nucleus is held together by the strong interaction, whereas a neutron star is held together by gravity. The density of a nucleus is uniform, while neutron stars are predicted to consist of multiple layers with varying compositions and densities.

Current constraints

Because equations of state for neutron stars lead to different observables, such as different mass-radius relations, there are many astronomical constraints on equations of state. These come mostly from LIGO, which is a gravitational wave observatory, and NICER, which is an X-ray telescope.

NICER's observations of pulsars in binary systems, from which the pulsar mass and radius can be estimated, can constrain the neutron star equation of state. A 2021 measurement of the pulsar PSR J0740+6620 was able to constrain the radius of a 1.4 solar mass neutron star to 12.33+0.76
−0.8 km with 95% confidence. These mass-radius constraints, combined with chiral effective field theory calculations, tightens constraints on the neutron star equation of state.

Equation of state constraints from LIGO gravitational wave detections start with nuclear and atomic physics researchers, who work to propose theoretical equations of state (such as FPS, UU, APR, L, SLy, and others). The proposed equations of state can then be passed onto astrophysics researchers who run simulations of binary neutron star mergers. From these simulations, researchers can extract gravitational waveforms, thus studying the relationship between the equation of state and gravitational waves emitted by binary neutron star mergers. Using these relations, one can constrain the neutron star equation of state when gravitational waves from binary neutron star mergers are observed. Past numerical relativity simulations of binary neutron star mergers have found relationships between the equation of state and frequency dependent peaks of the gravitational wave signal that can be applied to LIGO detections. For example, the LIGO detection of the binary neutron star merger GW170817 provided limits on the tidal deformability of the two neutron stars which dramatically reduced the family of allowed equations of state. Future gravitational wave signals with next generation detectors like Cosmic Explorer can impose further constraints.

When nuclear physicists are trying to understand the likelihood of their equation of state, it is good to compare with these constraints to see if it predicts neutron stars of these masses and radii. There is also recent work on constraining the equation of state with the speed of sound through hydrodynamics.

Tolman-Oppenheimer-Volkoff Equation

The Tolman-Oppenheimer-Volkoff (TOV) equation can be used to describe a neutron star. The equation is a solution to Einstein's equations from general relativity for a spherically symmetric, time invariant metric. With a given equation of state, solving the equation leads to observables such as the mass and radius. There are many codes that numerically solve the TOV equation for a given equation of state to find the mass-radius relation and other observables for that equation of state.

The following differential equations can be solved numerically to find the neutron star observables:

$${\displaystyle \frac{dp}{dr}=-\frac{Gϵ(r)M(r)}{c^2{r}^2}\left(1+\frac{p(r)}{ϵ(r)}\right)\left(1+\frac{4\pi r^{3}p(r)}{M(r)c^2}\right)\left(1-\frac{2GM(r)}{c^{2}r}\right)}$$ $${\displaystyle \frac{dM}{dr}=\frac{4\pi }{c^2}{r}^2ϵ(r)}$$

where ${\displaystyle G}$ is the gravitational constant, ${\displaystyle p(r)}$ is the pressure, ${\displaystyle ϵ(r)}$ is the energy density (found from the equation of state), and ${\displaystyle c}$ is the speed of light.

Mass-Radius relation

Using the TOV equations and an equation of state, a mass-radius curve can be found. The idea is that for the correct equation of state, every neutron star that could possibly exist would lie along that curve. This is one of the ways equations of state can be constrained by astronomical observations. To create these curves, one must solve the TOV equations for different central densities. For each central density, you numerically solve the mass and pressure equations until the pressure goes to zero, which is the outside of the star. Each solution gives a corresponding mass and radius for that central density.

Mass-radius curves determine what the maximum mass is for a given equation of state. Through most of the mass-radius curve, each radius corresponds to a unique mass value. At a certain point, the curve will reach a maximum and start going back down, leading to repeated mass values for different radii. This maximum point is what is known as the maximum mass. Beyond that mass, the star will no longer be stable, i.e. no longer be able to hold itself up against the force of gravity, and would collapse into a black hole. Since each equation of state leads to a different mass-radius curve, they also lead to a unique maximum mass value. The maximum mass value is unknown as long as the equation of state remains unknown.

This is very important when it comes to constraining the equation of state. Oppenheimer and Volkoff came up with the Tolman-Oppenheimer-Volkoff limit using a degenerate gas equation of state with the TOV equations that was ~0.7 Solar masses. Since the neutron stars that have been observed are more massive than that, that maximum mass was discarded. The most recent massive neutron star that was observed was PSR J0952-0607 which was 2.35±0.17 solar masses. Any equation of state with a mass less than that would not predict that star and thus is much less likely to be correct.

An interesting phenomenon in this area of astrophysics relating to the maximum mass of neutron stars is what is called the "mass gap". The mass gap refers to a range of masses from roughly 2-5 solar masses where very few compact objects were observed. This range is based on the current assumed maximum mass of neutron stars (~2 solar masses) and the minimum black hole mass (~5 solar masses). Recently, some objects have been discovered that fall in that mass gap from gravitational wave detections. If the true maximum mass of neutron stars was known, it would help characterize compact objects in that mass range as either neutron stars or black holes.

I-Love-Q Relations

There are three more properties of neutron stars that are dependent on the equation of state but can also be astronomically observed: the moment of inertia, the quadrupole moment, and the Love number. The moment of inertia of a neutron star describes how fast the star can rotate at a fixed spin momentum. The quadrupole moment of a neutron star specifies how much that star is deformed out of its spherical shape. The Love number of the neutron star represents how easy or difficult it is to deform the star due to tidal forces, typically important in binary systems.

While these properties depend on the material of the star and therefore on the equation of state, there is a relation between these three quantities that is independent of the equation of state. This relation assumes slowly and uniformly rotating stars and uses general relativity to derive the relation. While this relation would not be able to add constraints to the equation of state, since it is independent of the equation of state, it does have other applications. If one of these three quantities can be measured for a particular neutron star, this relation can be used to find the other two. In addition, this relation can be used to break the degeneracies in detections by gravitational wave detectors of the quadrupole moment and spin, allowing the average spin to be determined within a certain confidence level.

Temperature

The temperature inside a newly formed neutron star is from around 10¹¹ to 10¹² kelvin. However, the huge number of neutrinos it emits carries away so much energy that the temperature of an isolated neutron star falls within a few years to around 10⁶ kelvin. At this lower temperature, most of the light generated by a neutron star is in X-rays.

Some researchers have proposed a neutron star classification system using Roman numerals (not to be confused with the Yerkes luminosity classes for non-degenerate stars) to sort neutron stars by their mass and cooling rates: type I for neutron stars with low mass and cooling rates, type II for neutron stars with higher mass and cooling rates, and a proposed type III for neutron stars with even higher mass, approaching 2 M_☉, and with higher cooling rates and possibly candidates for exotic stars.

Magnetic field

The magnetic field strength on the surface of neutron stars ranges from c. 10⁴ to 10¹¹ tesla (T). These are orders of magnitude higher than in any other object: for comparison, a continuous 16 T field has been achieved in the laboratory and is sufficient to levitate a living frog due to diamagnetic levitation. Variations in magnetic field strengths are most likely the main factor that allows different types of neutron stars to be distinguished by their spectra, and explains the periodicity of pulsars.

The neutron stars known as magnetars have the strongest magnetic fields, in the range of 10⁸ to 10¹¹ T, and have become the widely accepted hypothesis for neutron star types soft gamma repeaters (SGRs) and anomalous X-ray pulsars (AXPs). The magnetic energy density of a 10⁸ T field is extreme, greatly exceeding the mass-energy density of ordinary matter. Fields of this strength are able to polarize the vacuum to the point that the vacuum becomes birefringent. Photons can merge or split in two, and virtual particle-antiparticle pairs are produced. The field changes electron energy levels and atoms are forced into thin cylinders. Unlike in an ordinary pulsar, magnetar spin-down can be directly powered by its magnetic field, and the magnetic field is strong enough to stress the crust to the point of fracture. Fractures of the crust cause starquakes, observed as extremely luminous millisecond hard gamma ray bursts. The fireball is trapped by the magnetic field, and comes in and out of view when the star rotates, which is observed as a periodic soft gamma repeater (SGR) emission with a period of 5–8 seconds and which lasts for a few minutes.

The origins of the strong magnetic field are as yet unclear. One hypothesis is that of "flux freezing", or conservation of the original magnetic flux during the formation of the neutron star. If an object has a certain magnetic flux over its surface area, and that area shrinks to a smaller area, but the magnetic flux is conserved, then the magnetic field would correspondingly increase. Likewise, a collapsing star begins with a much larger surface area than the resulting neutron star, and conservation of magnetic flux would result in a far stronger magnetic field. However, this simple explanation does not fully explain magnetic field strengths of neutron stars.

Gravity

See also: Tolman–Oppenheimer–Volkoff equation and White dwarf § Mass–radius relationship

Gravitational light deflection at a neutron star. Due to relativistic light deflection over half the surface is visible (each grid patch represents 30 by 30 degrees). In natural units, this star's mass is 1 and its radius is 4, or twice its Schwarzschild radius.

The gravitational field at a neutron star's surface is about 2×10¹¹ times stronger than on Earth, at around 2.0×10¹² m/s². Such a strong gravitational field acts as a gravitational lens and bends the radiation emitted by the neutron star such that parts of the normally invisible rear surface become visible. If the radius of the neutron star is 3GM/c² or less, then the photons may be trapped in an orbit, thus making the whole surface of that neutron star visible from a single vantage point, along with destabilizing photon orbits at or below the 1 radius distance of the star.

A fraction of the mass of a star that collapses to form a neutron star is released in the supernova explosion from which it forms (from the law of mass–energy equivalence, E = mc²). The energy comes from the gravitational binding energy of a neutron star.

Hence, the gravitational force of a typical neutron star is huge. If an object were to fall from a height of one meter on a neutron star 12 kilometers in radius, it would reach the ground at around 1,400 kilometers per second. However, even before impact, the tidal force would cause spaghettification, breaking any sort of an ordinary object into a stream of material.

Because of the enormous gravity, time dilation between a neutron star and Earth is significant. For example, eight years could pass on the surface of a neutron star, yet ten years would have passed on Earth, not including the time-dilation effect of the star's very rapid rotation.

Neutron star relativistic equations of state describe the relation of radius vs. mass for various models. The most likely radii for a given neutron star mass are bracketed by models AP4 (smallest radius) and MS2 (largest radius). E_B is the ratio of gravitational binding energy mass equivalent to the observed neutron star gravitational mass of M kilograms with radius R meters, $${\displaystyle E_{\text{B}}=\frac{0.60\beta }{1-\frac{\beta }{2}}}$$$${\displaystyle \beta =GM/R{c}^2}$$ Given current values

• ${\displaystyle G=6.67408\times {10}^{-11}{\text{m}}^3{\text{kg}}^{-1}{\text{s}}^{-2}}$
• ${\displaystyle c=2.99792458\times {10}^8\text{m}/\text{s}}$
• ${\displaystyle M_{\odot }=1.98855\times {10}^{30}\text{kg}}$

and star masses "M" commonly reported as multiples of one solar mass, $${\displaystyle M_x=\frac{M}{M_{\odot }}}$$ then the relativistic fractional binding energy of a neutron star is $${\displaystyle E_{\text{B}}=\frac{886.0M_x}{R_{\left[\text{in meters}\right]}-738.3M_x}}$$

A 2 M_☉ neutron star would not be more compact than 10,970 meters radius (AP4 model). Its mass fraction gravitational binding energy would then be 0.187, −18.7% (exothermic). This is not near 0.6/2 = 0.3, −30%.

Structure

Cross-section of neutron star. Densities are in terms of ρ₀ the saturation nuclear matter density, where nucleons begin to touch.

Current understanding of the structure of neutron stars is defined by existing mathematical models, but it might be possible to infer some details through studies of neutron-star oscillations. Asteroseismology, a study applied to ordinary stars, can reveal the inner structure of neutron stars by analyzing observed spectra of stellar oscillations.

Current models indicate that matter at the surface of a neutron star is composed of ordinary atomic nuclei crushed into a solid lattice with a sea of electrons flowing through the gaps between them. It is possible that the nuclei at the surface are iron, due to iron's high binding energy per nucleon. It is also possible that heavy elements, such as iron, simply sink beneath the surface, leaving only light nuclei like helium and hydrogen. If the surface temperature exceeds 10⁶ kelvins (as in the case of a young pulsar), the surface should be fluid instead of the solid phase that might exist in cooler neutron stars (temperature <10⁶ kelvins).

The "atmosphere" of a neutron star is hypothesized to be at most several micrometers thick, and its dynamics are fully controlled by the neutron star's magnetic field. Below the atmosphere one encounters a solid "crust". This crust is extremely hard and very smooth (with maximum surface irregularities on the order of millimeters or less), due to the extreme gravitational field.

Proceeding inward, one encounters nuclei with ever-increasing numbers of neutrons; such nuclei would decay quickly on Earth, but are kept stable by tremendous pressures. As this process continues at increasing depths, the neutron drip becomes overwhelming, and the concentration of free neutrons increases rapidly.

After a supernova explosion of a supergiant star, neutron stars are born from the remnants. A neutron star is composed mostly of neutrons (neutral particles) and contains a small fraction of protons (positively charged particles) and electrons (negatively charged particles), as well as nuclei. In the extreme density of a neutron star, many neutrons are free neutrons, meaning they are not bound in atomic nuclei and move freely within the star's dense matter, especially in the densest regions of the star—the inner crust and core. Over the star's lifetime, as its density increases, the energy of the electrons also increases, which generates more neutrons.

In neutron stars, the neutron drip is the transition point where nuclei become so neutron-rich that they can no longer hold additional neutrons, leading to a sea of free neutrons being formed. The sea of neutrons formed after neutron drip provides additional pressure support, which helps maintain the star's structural integrity and prevents gravitational collapse. The neutron drip takes place within the inner crust of the neutron star and starts when the density becomes so high that nuclei can no longer hold additional neutrons.

At the beginning of the neutron drip, the pressure in the star from neutrons, electrons, and the total pressure is roughly equal. As the density of the neutron star increases, the nuclei break down, and the neutron pressure of the star becomes dominant. When the density reaches a point where nuclei touch and subsequently merge, they form a fluid of neutrons with a sprinkle of electrons and protons. This transition marks the neutron drip, where the dominant pressure in the neutron star shifts from degenerate electrons to neutrons.

At very high densities, the neutron pressure becomes the primary pressure holding up the star, with neutrons being non-relativistic (moving slower than the speed of light) and extremely compressed. However, at extremely high densities, neutrons begin to move at relativistic speeds (close to the speed of light). These high speeds significantly increase the star's overall pressure, altering the star's equilibrium state, and potentially leading to the formation of exotic states of matter.

In that region, there are nuclei, free electrons, and free neutrons. The nuclei become increasingly small (gravity and pressure overwhelming the strong force) until the core is reached, by definition the point where mostly neutrons exist. The expected hierarchy of phases of nuclear matter in the inner crust has been characterized as "nuclear pasta", with fewer voids and larger structures towards higher pressures. The composition of the superdense matter in the core remains uncertain. One model describes the core as superfluid neutron-degenerate matter (mostly neutrons, with some protons and electrons). More exotic forms of matter are possible, including degenerate strange matter (containing strange quarks in addition to up and down quarks), matter containing high-energy pions and kaons in addition to neutrons, or ultra-dense quark-degenerate matter.

Radiation

Pulsars

Main article: Pulsar

Neutron stars are detected from their electromagnetic radiation. Neutron stars are usually observed to pulse radio waves and other electromagnetic radiation, and neutron stars observed with pulses are called pulsars.

Pulsars' radiation is thought to be caused by particle acceleration near their magnetic poles, which need not be aligned with the rotational axis of the neutron star. It is thought that a large electrostatic field builds up near the magnetic poles, leading to electron emission. These electrons are magnetically accelerated along the field lines, leading to curvature radiation, with the radiation being strongly polarized towards the plane of curvature. In addition, high-energy photons can interact with lower-energy photons and the magnetic field for electron−positron pair production, which through electron–positron annihilation leads to further high-energy photons.

The radiation emanating from the magnetic poles of neutron stars can be described as magnetospheric radiation, in reference to the magnetosphere of the neutron star. It is not to be confused with magnetic dipole radiation, which is emitted because the magnetic axis is not aligned with the rotational axis, with a radiation frequency the same as the neutron star's rotational frequency.

If the axis of rotation of the neutron star is different from the magnetic axis, external viewers will only see these beams of radiation whenever the magnetic axis point towards them during the neutron star rotation. Therefore, periodic pulses are observed, at the same rate as the rotation of the neutron star.

In May 2022, astronomers reported an ultra-long-period radio-emitting neutron star PSR J0901-4046, with spin properties distinct from the known neutron stars. It is unclear how its radio emission is generated, and it challenges the current understanding of how pulsars evolve.

Non-pulsating neutron stars

In addition to pulsars, non-pulsating neutron stars have also been identified, although they may have minor periodic variation in luminosity. This seems to be a characteristic of the X-ray sources known as Central Compact Objects in supernova remnants (CCOs in SNRs), which are thought to be young, radio-quiet isolated neutron stars.

Spectra

In addition to radio emissions, neutron stars have also been identified in other parts of the electromagnetic spectrum. This includes visible light, near infrared, ultraviolet, X-rays, and gamma rays. Pulsars observed in X-rays are known as X-ray pulsars if accretion-powered, while those identified in visible light are known as optical pulsars. The majority of neutron stars detected, including those identified in optical, X-ray, and gamma rays, also emit radio waves; the Crab Pulsar produces electromagnetic emissions across the spectrum. However, there exist neutron stars called radio-quiet neutron stars, with no radio emissions detected.

Rotation

Neutron stars rotate extremely rapidly after their formation due to the conservation of angular momentum; in analogy to spinning ice skaters pulling in their arms, the slow rotation of the original star's core speeds up as it shrinks. A newborn neutron star can rotate many times a second.

Spin down

P–P-dot diagram for known rotation-powered pulsars (red), anomalous X-ray pulsars (green), high-energy emission pulsars (blue) and binary pulsars (pink)

Over time, neutron stars slow, as their rotating magnetic fields in effect radiate energy associated with the rotation; older neutron stars may take several seconds for each revolution. This is called spin down. The rate at which a neutron star slows its rotation is usually constant and very small.

The periodic time (P) is the rotational period, the time for one rotation of a neutron star. The spin-down rate, the rate of slowing of rotation, is then given the symbol ${\displaystyle \dot{P}}$ (P-dot), the derivative of P with respect to time. It is defined as periodic time increase per unit time; it is a dimensionless quantity, but can be given the units of s⋅s⁻¹ (seconds per second).

The spin-down rate (P-dot) of neutron stars usually falls within the range of 10⁻²² to 10⁻⁹ s⋅s⁻¹, with the shorter period (or faster rotating) observable neutron stars usually having smaller P-dot. As a neutron star ages, its rotation slows (as P increases); eventually, the rate of rotation will become too slow to power the radio-emission mechanism, so radio emission from the neutron star no longer can be detected.

P and P-dot allow minimum magnetic fields of neutron stars to be estimated. P and P-dot can be also used to calculate the characteristic age of a pulsar, but gives an estimate which is somewhat larger than the true age when it is applied to young pulsars.

P and P-dot can also be combined with neutron star's moment of inertia to estimate a quantity called spin-down luminosity, which is given the symbol ${\displaystyle \dot{E}}$ (E-dot). It is not the measured luminosity, but rather the calculated loss rate of rotational energy that would manifest itself as radiation. For neutron stars where the spin-down luminosity is comparable to the actual luminosity, the neutron stars are said to be "rotation powered". The observed luminosity of the Crab Pulsar is comparable to the spin-down luminosity, supporting the model that rotational kinetic energy powers the radiation from it. With neutron stars such as magnetars, where the actual luminosity exceeds the spin-down luminosity by about a factor of one hundred, it is assumed that the luminosity is powered by magnetic dissipation, rather than being rotation powered.

P and P-dot can also be plotted for neutron stars to create a P–P-dot diagram. It encodes a tremendous amount of information about the pulsar population and its properties, and has been likened to the Hertzsprung–Russell diagram in its importance for neutron stars.

Spin up

A computer simulation depicting a neutron star with accretion disk, spewing out X-rays through the magnetic axis

Main article: Neutron star spin-up

Neutron star rotational speeds can increase, a process known as spin up. Sometimes neutron stars absorb orbiting matter from companion stars, increasing the rotation rate and reshaping the neutron star into an oblate spheroid. This causes an increase in the rate of rotation of the neutron star of over a hundred times per second in the case of millisecond pulsars.

The most rapidly rotating neutron star currently known, PSR J1748-2446ad, rotates at 716 revolutions per second. A 2007 paper reported the detection of an X-ray burst oscillation, which provides an indirect measure of spin, of 1122 Hz from the neutron star XTE J1739-285, suggesting 1122 rotations a second. However, at present, this signal has only been seen once, and should be regarded as tentative until confirmed in another burst from that star.

Glitches and starquakes

NASA artist's conception of a "starquake", or "stellar quake"

Sometimes a neutron star will undergo a glitch, a sudden small increase of its rotational speed or spin up. Glitches are thought to be the effect of a starquake—as the rotation of the neutron star slows, its shape becomes more spherical. Due to the stiffness of the "neutron" crust, this happens as discrete events when the crust ruptures, creating a starquake similar to earthquakes. After the starquake, the star will have a smaller equatorial radius, and because angular momentum is conserved, its rotational speed has increased.

Starquakes occurring in magnetars, with a resulting glitch, is the leading hypothesis for the gamma-ray sources known as soft gamma repeaters.

Recent work, however, suggests that a starquake would not release sufficient energy for a neutron star glitch; it has been suggested that glitches may instead be caused by transitions of vortices in the theoretical superfluid core of the neutron star from one metastable energy state to a lower one, thereby releasing energy that appears as an increase in the rotation rate.

Anti-glitches

An anti-glitch, a sudden small decrease in rotational speed, or spin down, of a neutron star has also been reported. It occurred in the magnetar 1E 2259+586, that in one case produced an X-ray luminosity increase of a factor of 20, and a significant spin-down rate change. Current neutron star models do not predict this behavior. If the cause were internal this suggests differential rotation of the solid outer crust and the superfluid component of the magnetar's inner structure.

Population and distances

At present, there are about 3,200 known neutron stars in the Milky Way and the Magellanic Clouds, the majority of which have been detected as radio pulsars. Neutron stars are mostly concentrated along the disk of the Milky Way, although the spread perpendicular to the disk is large because the supernova explosion process can impart high translational speeds (400 km/s) to the newly formed neutron star.

Some of the closest known neutron stars are RX J1856.5−3754, which is about 400 light-years from Earth, and PSR J0108−1431 about 424 light-years. RX J1856.5-3754 is a member of a close group of neutron stars called The Magnificent Seven. Another nearby neutron star that was detected transiting the backdrop of the constellation Ursa Minor has been nicknamed Calvera by its Canadian and American discoverers, after the villain in the 1960 film The Magnificent Seven. This rapidly moving object was discovered using the ROSAT Bright Source Catalog.

Neutron stars are only detectable with modern technology during the earliest stages of their lives (almost always less than 1 million years) and are vastly outnumbered by older neutron stars that would only be detectable through their blackbody radiation and gravitational effects on other stars.

Binary neutron star systems

Circinus X-1: X-ray light rings from a binary neutron star (24 June 2015; Chandra X-ray Observatory)

About 5% of all known neutron stars are members of a binary system. The formation and evolution of binary neutron stars and double neutron stars can be a complex process. Neutron stars have been observed in binaries with ordinary main-sequence stars, red giants, white dwarfs, or other neutron stars. According to modern theories of binary evolution, it is expected that neutron stars also exist in binary systems with black hole companions. The merger of binaries containing two neutron stars, or a neutron star and a black hole, has been observed through the emission of gravitational waves.

X-ray binaries

Main article: X-ray binary

Binary systems containing neutron stars often emit X-rays, which are emitted by hot gas as it falls towards the surface of the neutron star. The source of the gas is the companion star, the outer layers of which can be stripped off by the gravitational force of the neutron star if the two stars are sufficiently close. As the neutron star accretes this gas, its mass can increase; if enough mass is accreted, the neutron star may collapse into a black hole.

Neutron star binary mergers and nucleosynthesis

Main article: Neutron star merger

 

Four snapshots from a computer simulation of a neutron star merger. 

1. The two neutron stars make initial contact
2. Immense tidal forces begin to disrupt the outer layers of the neutron stars
3. The neutron stars are completely tidally disrupted
4. A black hole forms, surrounded by an accretion disc

The distance between two neutron stars in a close binary system is observed to shrink as gravitational waves are emitted. Ultimately, the neutron stars will come into contact and coalesce. The coalescence of binary neutron stars is one of the leading models for the origin of short gamma-ray bursts. Strong evidence for this model came from the observation of a kilonova associated with the short-duration gamma-ray burst GRB 130603B, and was finally confirmed by detection of gravitational wave GW170817 and short GRB 170817A by LIGO, Virgo, and 70 observatories covering the electromagnetic spectrum observing the event. The light emitted in the kilonova is believed to come from the radioactive decay of material ejected in the merger of the two neutron stars. The merger momentarily creates an environment of such extreme neutron flux that the r-process can occur; this—as opposed to supernova nucleosynthesis—may be responsible for the production of around half the isotopes in chemical elements beyond iron.

Planets

Main articles: Pulsar planet and Habitability of neutron star systems

Neutron stars can host exoplanets. These can be original, circumbinary, captured, or the result of a second round of planet formation. Pulsars can also strip the atmosphere off from a star, leaving a planetary-mass remnant, which may be understood as a chthonian planet or a stellar object depending on interpretation. For pulsars, such pulsar planets can be detected with the pulsar timing method, which allows for high precision and detection of much smaller planets than with other methods. Two systems have been definitively confirmed. The first exoplanets ever to be detected were the three planets Draugr, Poltergeist and Phobetor around the pulsar Lich, discovered in 1992–1994. Of these, Draugr is the smallest exoplanet ever detected, at a mass of twice that of the Moon. Another system is PSR B1620−26, where a circumbinary planet orbits a neutron star-white dwarf binary system. Also, there are several unconfirmed candidates. Pulsar planets receive little visible light, but massive amounts of ionizing radiation and high-energy stellar wind, which makes them rather hostile environments to life as presently understood.

History of discoveries

The first direct observation of an isolated neutron star in visible light. The neutron star is RX J1856.5−3754.

At the meeting of the American Physical Society in December 1933 (the proceedings were published in January 1934), Walter Baade and Fritz Zwicky proposed the existence of neutron stars, less than two years after the discovery of the neutron by James Chadwick. In seeking an explanation for the origin of a supernova, they tentatively proposed that in supernova explosions ordinary stars are turned into stars that consist of extremely closely packed neutrons that they called neutron stars. Baade and Zwicky correctly proposed at that time that the release of the gravitational binding energy of the neutron stars powers the supernova: "In the supernova process, mass in bulk is annihilated". Neutron stars were thought to be too faint to be detectable and little work was done on them until November 1967, when Franco Pacini pointed out that if the neutron stars were spinning and had large magnetic fields, then electromagnetic waves would be emitted. Unknown to him, radio astronomer Antony Hewish and his graduate student Jocelyn Bell at Cambridge were shortly to detect radio pulses from stars that are now believed to be highly magnetized, rapidly spinning neutron stars, known as pulsars.

In 1965, Antony Hewish and Samuel Okoye discovered "an unusual source of high radio brightness temperature in the Crab Nebula". This source turned out to be the Crab Pulsar that resulted from the great supernova of 1054.

In 1967, Iosif Shklovsky examined the X-ray and optical observations of Scorpius X-1 and correctly concluded that the radiation comes from a neutron star at the stage of accretion.

In 1967, Jocelyn Bell Burnell and Antony Hewish discovered regular radio pulses from PSR B1919+21. This pulsar was later interpreted as an isolated, rotating neutron star. The energy source of the pulsar is the rotational energy of the neutron star. The majority of known neutron stars (about 2000, as of 2010) have been discovered as pulsars, emitting regular radio pulses.

In 1968, Richard V. E. Lovelace and collaborators discovered period ${\displaystyle P\approx 33}$ ms of the Crab pulsar using Arecibo Observatory. After this discovery, scientists concluded that pulsars were rotating neutron stars. Before that, many scientists believed that pulsars were pulsating white dwarfs.

In 1971, Riccardo Giacconi, Herbert Gursky, Ed Kellogg, R. Levinson, E. Schreier, and H. Tananbaum discovered 4.8 second pulsations in an X-ray source in the constellation Centaurus, Cen X-3. They interpreted this as resulting from a rotating hot neutron star. The energy source is gravitational and results from a rain of gas falling onto the surface of the neutron star from a companion star or the interstellar medium.

In 1974, Antony Hewish was awarded the Nobel Prize in Physics "for his decisive role in the discovery of pulsars" without Jocelyn Bell who shared in the discovery.

In 1974, Joseph Taylor and Russell Hulse discovered the first binary pulsar, PSR B1913+16, which consists of two neutron stars (one seen as a pulsar) orbiting around their center of mass. Albert Einstein's general theory of relativity predicts that massive objects in short binary orbits should emit gravitational waves, and thus that their orbit should decay with time. This was indeed observed, precisely as general relativity predicts, and in 1993, Taylor and Hulse were awarded the Nobel Prize in Physics for this discovery.

In 1982, Don Backer and colleagues discovered the first millisecond pulsar, PSR B1937+21. This object spins 642 times per second, a value that placed fundamental constraints on the mass and radius of neutron stars. Many millisecond pulsars were later discovered, but PSR B1937+21 remained the fastest-spinning known pulsar for 24 years, until PSR J1748-2446ad (which spins ~716 times a second) was discovered.

In 2003, Marta Burgay and colleagues discovered the first double neutron star system where both components are detectable as pulsars, PSR J0737−3039. The discovery of this system allows a total of 5 different tests of general relativity, some of these with unprecedented precision.

In 2010, Paul Demorest and colleagues measured the mass of the millisecond pulsar PSR J1614−2230 to be 1.97±0.04 M_☉, using Shapiro delay. This was substantially higher than any previously measured neutron star mass (1.67 M_☉, see PSR J1903+0327), and places strong constraints on the interior composition of neutron stars.

In 2013, John Antoniadis and colleagues measured the mass of PSR J0348+0432 to be 2.01±0.04 M_☉, using white dwarf spectroscopy. This confirmed the existence of such massive stars using a different method. Furthermore, this allowed, for the first time, a test of general relativity using such a massive neutron star.

In August 2017, LIGO and Virgo made first detection of gravitational waves produced by colliding neutron stars (GW170817), leading to further discoveries about neutron stars.

In October 2018, astronomers reported that GRB 150101B, a gamma-ray burst event detected in 2015, may be directly related to the historic GW170817 and associated with the merger of two neutron stars. The similarities between the two events, in terms of gamma ray, optical and x-ray emissions, as well as to the nature of the associated host galaxies, are "striking", suggesting the two separate events may both be the result of the merger of neutron stars, and both may be a kilonova, which may be more common in the universe than previously understood, according to the researchers.

In July 2019, astronomers reported that a new method to determine the Hubble constant, and resolve the discrepancy of earlier methods, has been proposed based on the mergers of pairs of neutron stars, following the detection of the neutron star merger of GW170817. Their measurement of the Hubble constant is 70.3+5.3
−5.0 (km/s)/Mpc.

A 2020 study by University of Southampton PhD student Fabian Gittins suggested that surface irregularities ("mountains") may only be fractions of a millimeter tall (about 0.000003% of the neutron star's diameter), hundreds of times smaller than previously predicted, a result bearing implications for the non-detection of gravitational waves from spinning neutron stars.

Using the JWST, astronomers have identified a neutron star within the remnants of the Supernova 1987A stellar explosion after seeking to do so for 37 years, according to a 23 February 2024 Science article. In a paradigm shift, new JWST data provides the elusive direct confirmation of neutron stars within supernova remnants as well as a deeper understanding of the processes at play within SN 1987A's remnants.

Subtypes

Different Types of Neutron Stars

Computer renders of a neutron star with accretion disk, with magnetic field lines projected, showing bursts of powerful X-rays. The simulations are taken from 2017 data from NASA's NuSTAR and Swift, and ESA's XMM-Newton observatories.

There are a number of types of object that consist of or contain a neutron star:

• Isolated neutron star (INS): not in a binary system.
  • Rotation-powered pulsar (RPP or "radio pulsar"): neutron stars that emit directed pulses of radiation towards us at regular intervals (due to their strong magnetic fields).
    • Rotating radio transient (RRATs): are thought to be pulsars which emit more sporadically and/or with higher pulse-to-pulse variability than the bulk of the known pulsars.
  • Magnetar: a neutron star with an extremely strong magnetic field (1000 times more than a regular neutron star), and long rotation periods (5 to 12 seconds).
    • Soft gamma repeater (SGR).
    • Anomalous X-ray pulsar (AXP).
  • Radio-quiet neutron stars.
    • X-ray dim isolated neutron stars.
    • Central compact objects in supernova remnants (CCOs in SNRs): young, radio-quiet non-pulsating X-ray sources, thought to be Isolated Neutron Stars surrounded by supernova remnants.
• X-ray pulsars or "accretion-powered pulsars": a class of X-ray binaries.
  • Low-mass X-ray binary pulsars: a class of low-mass X-ray binaries (LMXB), a pulsar with a main sequence star, white dwarf or red giant.
    • Millisecond pulsar (MSP) ("recycled pulsar").
      • "Spider Pulsar", a pulsar where their companion is a semi-degenerate star.
        • "Black Widow" pulsar, a pulsar that falls under the "Spider Pulsar" if the companion has extremely low mass (less than 0.1 M_☉).
        • "Redback" pulsar, are if the companion is more massive.
      • Sub-millisecond pulsar.
    • X-ray burster: a neutron star with a low mass binary companion from which matter is accreted resulting in irregular bursts of energy from the surface of the neutron star.
  • Intermediate-mass X-ray binary pulsars: a class of intermediate-mass X-ray binaries (IMXB), a pulsar with an intermediate mass star.
  • High-mass X-ray binary pulsars: a class of high-mass X-ray binaries (HMXB), a pulsar with a massive star.
  • Binary pulsars: a pulsar with a binary companion, often a white dwarf or neutron star.
  • X-ray tertiary (theorized).

There are also a number of theorized compact stars with similar properties that are not actually neutron stars.

• Protoneutron star (PNS), a theorized intermediate–stage object that cools and contracts to form a neutron star or a black hole
• Exotic star
  • Thorne–Żytkow object: currently a hypothetical merger of a neutron star into a red giant star.
  • Quark star: currently a hypothetical type of neutron star composed of quark matter, or strange matter. As of 2018, there are three candidates.
  • Electroweak star: currently a hypothetical type of extremely heavy neutron star, in which the quarks are converted to leptons through the electroweak force, but the gravitational collapse of the neutron star is prevented by radiation pressure. As of 2018, there is no evidence for their existence.
  • Preon star: currently a hypothetical type of neutron star composed of preon matter. As of 2018, there is no evidence for the existence of preons.

Examples of neutron stars

An artist's conception of the pulsar planet PSR B1257+12 C, with bright aurorae

See also: Pulsar § Significant pulsars

• Black Widow Pulsar – a millisecond pulsar that is very massive
• PSR J0952-0607 – the heaviest neutron star with 2.35+0.17
  −0.17 M_☉, a type of Black Widow Pulsar
• LGM-1 (now known as PSR B1919+21) – the first recognized radio-pulsar. It was discovered by Jocelyn Bell Burnell in 1967.
• PSR B1257+12 (Also known as Lich) – the first neutron star discovered with planets (a millisecond pulsar).
• PSR B1509−58 – source of the "Hand of God" photo shot by the Chandra X-ray Observatory
• RX J1856.5−3754 – the closest neutron star
• The Magnificent Seven – a group of nearby, X-ray dim isolated neutron stars
• PSR J0348+0432 – the most massive neutron star with a well-constrained mass, 2.01±0.04 M_☉
• SWIFT J1756.9-2508 – a millisecond pulsar with a stellar-type companion with planetary range mass (below brown dwarf)
• Swift J1818.0-1607 – the youngest-known magnetar

Globular cluster

From Wikipedia, the free encyclopedia

https://en.wikipedia.org/wiki/Globular_cluster 

 

Globular cluster
Messier 2
Characteristics
Type	Star cluster
Mass range	1K M☉ - >1M M☉
Size range	10-300 ly across
Density	~2 stars/cubic ly
Average luminosity	~25 000 L☉

A globular cluster is a spheroidal conglomeration of stars that is bound together by gravity, with a higher concentration of stars towards its center. It can contain anywhere from tens of thousands to many millions of member stars, all orbiting in a stable, compact formation. Globular clusters are similar in form to dwarf spheroidal galaxies, and though globular clusters were long held to be the more luminous of the two, discoveries of outliers had made the distinction between the two less clear by the early 21st century. Their name is derived from Latin globulus (small sphere). Globular clusters are occasionally known simply as "globulars".

Although one globular cluster, Omega Centauri, was observed in antiquity and long thought to be a star, recognition of the clusters' true nature came with the advent of telescopes in the 17th century. In early telescopic observations, globular clusters appeared as fuzzy blobs, leading French astronomer Charles Messier to include many of them in his catalog of astronomical objects that he thought could be mistaken for comets. Using larger telescopes, 18th-century astronomers recognized that globular clusters are groups of many individual stars. Early in the 20th century the distribution of globular clusters in the sky was some of the first evidence that the Sun is far from the center of the Milky Way.

Globular clusters are found in nearly all galaxies. In spiral galaxies like the Milky Way, they are mostly found in the outer spheroidal part of the galaxy – the galactic halo. They are the largest and most massive type of star cluster, tending to be older, denser, and composed of lower abundances of heavy elements than open clusters, which are generally found in the disks of spiral galaxies. The Milky Way has more than 150 known globulars, and there may be many more.

Both the origin of globular clusters and their role in galactic evolution are unclear. Some are among the oldest objects in their galaxies and even the universe, constraining estimates of the universe's age. Star clusters were formerly thought to consist of stars that all formed at the same time from one star-forming nebula, but nearly all globular clusters contain stars that formed at different times, or that have differing compositions. Some clusters may have had multiple episodes of star formation, and some may be remnants of smaller galaxies captured by larger galaxies.

History of observations

The first known globular cluster, now called M 22, was discovered in 1665 by Abraham Ihle, a German amateur astronomer. The cluster Omega Centauri, easily visible in the southern sky with the naked eye, was known to ancient astronomers like Ptolemy as a star, but was reclassified as a nebula by Edmond Halley in 1677, then finally as a globular cluster in the early 19th century by John Herschel. The French astronomer Abbé Lacaille listed NGC 104, NGC 4833, M 55, M 69, and NGC 6397 in his 1751–1752 catalogue.^[a] The low resolution of early telescopes prevented individual stars in a cluster from being visually separated until Charles Messier observed M 4 in 1764.

Early globular cluster discoveries

Cluster name	Discovered by	Year
M 22	Abraham Ihle	1665
ω Cen	Edmond Halley	1677
M 5	Gottfried Kirch	1702
M 13	Edmond Halley	1714
M 71	Philippe Loys de Chéseaux	1745
M 4	Philippe Loys de Chéseaux	1746
M 15	Jean-Dominique Maraldi	1746
M 2	Jean-Dominique Maraldi	1746

When William Herschel began his comprehensive survey of the sky using large telescopes in 1782, there were 34 known globular clusters. Herschel discovered another 36 and was the first to resolve virtually all of them into stars. He coined the term globular cluster in his Catalogue of a Second Thousand New Nebulae and Clusters of Stars (1789). In 1914, Harlow Shapley began a series of studies of globular clusters, published across about forty scientific papers. He examined the clusters' RR Lyrae variables (stars which he assumed were Cepheid variables) and used their luminosity and period of variability to estimate the distances to the clusters. RR Lyrae variables were later found to be fainter than Cepheid variables, causing Shapley to overestimate the distances.

NGC 7006 is a highly concentrated, Class I globular cluster.

A large majority of the Milky Way's globular clusters are found in the halo around the galactic core. In 1918, Shapley used this strongly asymmetrical distribution to determine the overall dimensions of the galaxy. Assuming a roughly spherical distribution of globular clusters around the galaxy's center, he used the positions of the clusters to estimate the position of the Sun relative to the Galactic Center. He correctly concluded that the Milky Way's center is in the Sagittarius constellation and not near the Earth. He overestimated the distance, finding typical globular cluster distances of 10–30 kiloparsecs (33,000–98,000 ly); the modern distance to the Galactic Center is roughly 8.5 kiloparsecs (28,000 ly). Shapley's measurements indicated the Sun is relatively far from the center of the galaxy, contrary to what had been inferred from the observed uniform distribution of ordinary stars. In reality most ordinary stars lie within the galaxy's disk and are thus obscured by gas and dust in the disk, whereas globular clusters lie outside the disk and can be seen at much greater distances.

The Messier 80 globular cluster in the constellation Scorpius is located about 30,000 light-years from the Sun and contains hundreds of thousands of stars.

The count of known globular clusters in the Milky Way has continued to increase, reaching 83 in 1915, 93 in 1930, 97 by 1947, and 157 in 2010. The number of known globular clusters in the Milky Way reached 158 by the end of 2010, according to the European Southern Observatory, before two new globular clusters were discovered as part of the ESO’s VISTA (Visible and Infrared Survey Telescope for Astronomy) infrared survey, known as Variables in the Vía Láctea (VVV) survey, bringing the total to 160 known globular clusters. The two discovered by VISTA in 2011 are named VVV CL001 and VVV CL002.

Additional, undiscovered globular clusters are believed to be in the galactic bulge or hidden by the gas and dust of the Milky Way. For example, most of the Palomar Globular Clusters have only been discovered in the 1950s, with some located relatively close-by yet obscured by dust, while others reside in the very far reaches of the Milky Way halo. The Andromeda Galaxy, which is comparable in size to the Milky Way, may have as many as five hundred globulars. Every galaxy of sufficient mass in the Local Group has an associated system of globular clusters, as does almost every large galaxy surveyed. Some giant elliptical galaxies (particularly those at the centers of galaxy clusters), such as M 87, have as many as 13,000 globular clusters.

Classification

Main article: Shapley–Sawyer Concentration Class

Shapley was later assisted in his studies of clusters by Henrietta Swope and Helen Sawyer Hogg. In 1927–1929, Shapley and Sawyer categorized clusters by the degree of concentration of stars toward each core. Their system, known as the Shapley–Sawyer Concentration Class, identifies the most concentrated clusters as Class I and ranges to the most diffuse Class XII. Astronomers from the Pontifical Catholic University of Chile proposed a new type of globular cluster on the basis of observational data in 2015: Dark globular clusters.

Formation

NGC 2808 contains three distinct generations of stars.
NASA image

The formation of globular clusters is poorly understood. Globular clusters have traditionally been described as a simple star population formed from a single giant molecular cloud, and thus with roughly uniform age and metallicity (proportion of heavy elements in their composition). Modern observations show that nearly all globular clusters contain multiple populations; the globular clusters in the Large Magellanic Cloud (LMC) exhibit a bimodal population, for example. During their youth, these LMC clusters may have encountered giant molecular clouds that triggered a second round of star formation. This star-forming period is relatively brief, compared with the age of many globular clusters. It has been proposed that this multiplicity in stellar populations could have a dynamical origin. In the Antennae Galaxy, for example, the Hubble Space Telescope has observed clusters of clusters – regions in the galaxy that span hundreds of parsecs, in which many of the clusters will eventually collide and merge. Their overall range of ages and (possibly) metallicities could lead to clusters with a bimodal, or even multiple, distribution of populations.

Globular star cluster Messier 54

Observations of globular clusters show that their stars primarily come from regions of more efficient star formation, and from where the interstellar medium is at a higher density, as compared to normal star-forming regions. Globular cluster formation is prevalent in starburst regions and in interacting galaxies. Some globular clusters likely formed in dwarf galaxies and were removed by tidal forces to join the Milky Way. In elliptical and lenticular galaxies there is a correlation between the mass of the supermassive black holes (SMBHs) at their centers and the extent of their globular cluster systems. The mass of the SMBH in such a galaxy is often close to the combined mass of the galaxy's globular clusters.

No known globular clusters display active star formation, consistent with the hypothesis that globular clusters are typically the oldest objects in their galaxy and were among the first collections of stars to form. Very large regions of star formation known as super star clusters, such as Westerlund 1 in the Milky Way, may be the precursors of globular clusters.

Many of the Milky Way's globular clusters have a retrograde orbit (meaning that they revolve around the galaxy in the reverse of the direction the galaxy is rotating), including the most massive, Omega Centauri. Its retrograde orbit suggests it may be a remnant of a dwarf galaxy captured by the Milky Way.

Composition

Djorgovski 1's stars contain hydrogen and helium, but not much else. In astronomical terms they are metal-poor.

Globular clusters are generally composed of hundreds of thousands of low-metal, old stars. The stars found in a globular cluster are similar to those in the bulge of a spiral galaxy but confined to a spheroid in which half the light is emitted within a radius of only a few to a few tens of parsecs. They are free of gas and dust, and it is presumed that all the gas and dust was long ago either turned into stars or blown out of the cluster by the massive first-generation stars.

Globular clusters can contain a high density of stars; on average about 0.4 stars per cubic parsec, increasing to 100 or 1000 stars/pc³ in the core of the cluster. In comparison, the stellar density around the Sun is roughly 0.1 stars/pc³. The typical distance between stars in a globular cluster is about one light year, but at its core the separation between stars averages about a third of a light year – thirteen times closer than the Sun is to its nearest neighbor, Proxima Centauri.

Globular clusters are thought to be unfavorable locations for planetary systems. Planetary orbits are dynamically unstable within the cores of dense clusters because of the gravitational perturbations of passing stars. A planet orbiting at one astronomical unit around a star that is within the core of a dense cluster, such as 47 Tucanae, would survive only on the order of a hundred million years. There is a planetary system orbiting a pulsar (PSR B1620−26) that belongs to the globular cluster M4, but these planets likely formed after the event that created the pulsar.

Some globular clusters, like Omega Centauri in the Milky Way and Mayall II in the Andromeda Galaxy, are extraordinarily massive, measuring several million solar masses (M_☉) and having multiple stellar populations. Both are evidence that supermassive globular clusters formed from the cores of dwarf galaxies that have been consumed by larger galaxies. About a quarter of the globular cluster population in the Milky Way may have been accreted this way, as with more than 60% of the globular clusters in the outer halo of Andromeda.

Heavy element content

Globular clusters normally consist of Population II stars which, compared with Population I stars such as the Sun, have a higher proportion of hydrogen and helium and a lower proportion of heavier elements. Astronomers refer to these heavier elements as metals (distinct from the material concept) and to the proportions of these elements as the metallicity. Produced by stellar nucleosynthesis, the metals are recycled into the interstellar medium and enter a new generation of stars. The proportion of metals can thus be an indication of the age of a star in simple models, with older stars typically having a lower metallicity.

The Dutch astronomer Pieter Oosterhoff observed two special populations of globular clusters, which became known as Oosterhoff groups. The second group has a slightly longer period of RR Lyrae variable stars. While both groups have a low proportion of metallic elements as measured by spectroscopy, the metal spectral lines in the stars of Oosterhoff type I (Oo I) cluster are not quite as weak as those in type II (Oo II), and so type I stars are referred to as metal-rich (e.g. Terzan 7), while type II stars are metal-poor (e.g. ESO 280-SC06). These two distinct populations have been observed in many galaxies, especially massive elliptical galaxies. Both groups are nearly as old as the universe itself and are of similar ages. Suggested scenarios to explain these subpopulations include violent gas-rich galaxy mergers, the accretion of dwarf galaxies, and multiple phases of star formation in a single galaxy. In the Milky Way, the metal-poor clusters are associated with the halo and the metal-rich clusters with the bulge.

A large majority of the metal-poor clusters in the Milky Way are aligned on a plane in the outer part of the galaxy's halo. This observation supports the view that type II clusters were captured from a satellite galaxy, rather than being the oldest members of the Milky Way's globular cluster system as was previously thought. The difference between the two cluster types would then be explained by a time delay between when the two galaxies formed their cluster systems.

Exotic components

Messier 53 contains an unusually large number of a type of star called blue stragglers.

Close interactions and near-collisions of stars occur relatively often in globular clusters because of their high star density. These chance encounters give rise to some exotic classes of stars – such as blue stragglers, millisecond pulsars, and low-mass X-ray binaries – which are much more common in globular clusters. How blue stragglers form remains unclear, but most models attribute them to interactions between stars, such as stellar mergers, the transfer of material from one star to another, or even an encounter between two binary systems. The resulting star has a higher temperature than other stars in the cluster with comparable luminosity and thus differs from the main-sequence stars formed early in the cluster's existence. Some clusters have two distinct sequences of blue stragglers, one bluer than the other.

Globular cluster M15 may have an intermediate-mass black hole at its core, but this claim is contested.

Simulation of stellar motions in Messier 4, where astronomers suspect that an intermediate-mass black hole could be present. If confirmed, the black hole would be in the center of the cluster, and would have a sphere of influence (black hole) limited by the red circle.

Astronomers have searched for black holes within globular clusters since the 1970s. The required resolution for this task is exacting; it is only with the Hubble Space Telescope (HST) that the first claimed discoveries were made, in 2002 and 2003. Based on HST observations, other researchers suggested the existence of a 4,000 M_☉(solar masses) intermediate-mass black hole in the globular cluster M15 and a 20,000 M_☉ black hole in the Mayall II cluster of the Andromeda Galaxy. Both X-ray and radio emissions from Mayall II appear consistent with an intermediate-mass black hole; however, these claimed detections are controversial.

The heaviest objects in globular clusters are expected to migrate to the cluster center due to mass segregation. One research group pointed out that the mass-to-light ratio should rise sharply towards the center of the cluster, even without a black hole, in both M15 and Mayall II. Observations from 2018 find no evidence for an intermediate-mass black hole in any globular cluster, including M15, but cannot definitively rule out one with a mass of 500–1000 M_☉. Finally, in 2023, an analysis of HST and the Gaia spacecraft data from the closest globular cluster, Messier 4, revealed an excess mass of roughly 800 M_☉ in the center of this cluster, which appears to not be extended. This could thus be considered as kinematic evidence for an intermediate-mass black hole (even if an unusually compact cluster of compact objects like white dwarfs, neutron stars or stellar-mass black holes cannot be completely discounted).

The confirmation of intermediate-mass black holes in globular clusters would have important ramifications for theories of galaxy development as being possible sources for the supermassive black holes at their centers. The mass of these supposed intermediate-mass black holes is proportional to the mass of their surrounding clusters, following a pattern previously discovered between supermassive black holes and their surrounding galaxies.

Hertzsprung–Russell diagrams

H–R diagram for the globular cluster M3. There is a characteristic "knee" in the curve at magnitude 19 where stars begin entering the giant stage of their evolutionary path, the main-sequence turnoff.

Hertzsprung–Russell diagrams (H–R diagrams) of globular clusters allow astronomers to determine many of the properties of their populations of stars. An H–R diagram is a graph of a large sample of stars plotting their absolute magnitude (their luminosity, or brightness measured from a standard distance), as a function of their color index. The color index, roughly speaking, measures the color of the star; positive color indices indicate a reddish star with a cool surface temperature, while negative values indicate a bluer star with a hotter surface. Stars on an H–R diagram mostly lie along a roughly diagonal line sloping from hot, luminous stars in the upper left to cool, faint stars in the lower right. This line is known as the main sequence and represents the primary stage of stellar evolution. The diagram also includes stars in later evolutionary stages such as the cool but luminous red giants.

Constructing an H–R diagram requires knowing the distance to the observed stars to convert apparent into absolute magnitude. Because all the stars in a globular cluster have about the same distance from Earth, a color–magnitude diagram using their observed magnitudes looks like a shifted H–R diagram (because of the roughly constant difference between their apparent and absolute magnitudes). This shift is called the distance modulus and can be used to calculate the distance to the cluster. The modulus is determined by comparing features (like the main sequence) of the cluster's color–magnitude diagram to corresponding features in an H–R diagram of another set of stars, a method known as spectroscopic parallax or main-sequence fitting.

Properties

Since globular clusters form at once from a single giant molecular cloud, a cluster's stars have roughly the same age and composition. A star's evolution is primarily determined by its initial mass, so the positions of stars in a cluster's H–R or color–magnitude diagram mostly reflect their initial masses. A cluster's H–R diagram, therefore, appears quite different from H–R diagrams containing stars of a wide variety of ages. Almost all stars fall on a well-defined curve in globular cluster H–R diagrams, and that curve's shape indicates the age of the cluster. A more detailed H–R diagram often reveals multiple stellar populations as indicated by the presence of closely separated curves, each corresponding to a distinct population of stars with a slightly different age or composition. Observations with the Wide Field Camera 3, installed in 2009 on the Hubble Space Telescope, made it possible to distinguish these slightly different curves.

The most massive main-sequence stars have the highest luminosity and will be the first to evolve into the giant star stage. As the cluster ages, stars of successively lower masses will do the same. Therefore, the age of a single-population cluster can be measured by looking for those stars just beginning to enter the giant star stage, which form a "knee" in the H–R diagram called the main-sequence turnoff, bending to the upper right from the main-sequence line. The absolute magnitude at this bend is directly a function of the cluster's age; an age scale can be plotted on an axis parallel to the magnitude.

The morphology and luminosity of globular cluster stars in H–R diagrams are influenced by numerous parameters, many of which are still actively researched. Recent observations have overturned the historical paradigm that all globular clusters consist of stars born at exactly the same time, or sharing exactly the same chemical abundance. Some clusters feature multiple populations, slightly differing in composition and age; for example, high-precision imagery of cluster NGC 2808 discerned three close, but distinct, main sequences. Further, the placements of the cluster stars in an H–R diagram (including the brightnesses of distance indicators) can be influenced by observational biases. One such effect, called blending, arises when the cores of globular clusters are so dense that observations see multiple stars as a single target. The brightness measured for that seemingly single star is thus incorrect – too bright, given that multiple stars contributed. In turn, the computed distance is incorrect, so the blending effect can introduce a systematic uncertainty into the cosmic distance ladder and may bias the estimated age of the universe and the Hubble constant.

Consequences

The blue stragglers appear on the H–R diagram as a series diverging from the main sequence in the direction of brighter, bluer stars. White dwarfs (the final remnants of some Sun-like stars), which are much fainter and somewhat hotter than the main-sequence stars, lie on the bottom-left of an H–R diagram. Globular clusters can be dated by looking at the temperatures of the coolest white dwarfs, often giving results as old as 12.7 billion years. In comparison, open clusters are rarely older than about half a billion years. The ages of globular clusters place a lower bound on the age of the entire universe, presenting a significant constraint in cosmology. Astronomers were historically faced with age estimates of clusters older than their cosmological models would allow, but better measurements of cosmological parameters, through deep sky surveys and satellites, appear to have resolved this issue.

Studying globular clusters sheds light on how the composition of the formational gas and dust affects stellar evolution; the stars' evolutionary tracks vary depending on the abundance of heavy elements. Data obtained from these studies are then used to study the evolution of the Milky Way as a whole.

Morphology

Ellipticity of globular clusters

Galaxy	Ellipticity
Milky Way	0.07±0.04
LMC	0.16±0.05
SMC	0.19±0.06
M31	0.09±0.04

In contrast to open clusters, most globular clusters remain gravitationally bound together for time periods comparable to the lifespans of most of their stars. Strong tidal interactions with other large masses result in the dispersal of some stars, leaving behind "tidal tails" of stars removed from the cluster.

After formation, the stars in the globular cluster begin to interact gravitationally with each other. The velocities of the stars steadily change, and the stars lose any history of their original velocity. The characteristic interval for this to occur is the relaxation time, related to the characteristic length of time a star needs to cross the cluster and the number of stellar masses. The relaxation time varies by cluster, but a typical value is on the order of one billion years.

Although globular clusters are generally spherical in form, ellipticity can form via tidal interactions. Clusters within the Milky Way and the Andromeda Galaxy are typically oblate spheroids in shape, while those in the Large Magellanic Cloud are more elliptical.

Radii

"Tidal radius" redirects here and is not to be confused with Roche limit.

NGC 411 is classified as an open cluster.

Astronomers characterize the morphology (shape) of a globular cluster by means of standard radii: the core radius (r_c), the half-light radius (r_h), and the tidal or Jacobi radius (r_t). The radius can be expressed as a physical distance or as a subtended angle in the sky. Considering a radius around the core, the surface luminosity of the cluster steadily decreases with distance, and the core radius is the distance at which the apparent surface luminosity has dropped by half. A comparable quantity is the half-light radius, or the distance from the core containing half the total luminosity of the cluster; the half-light radius is typically larger than the core radius.

Most globular clusters have a half-light radius of less than ten parsecs (pc), although some globular clusters have very large radii, like NGC 2419 (r_h = 18 pc) and Palomar 14 (r_h = 25 pc). The half-light radius includes stars in the outer part of the cluster that happen to lie along the line of sight, so theorists also use the half-mass radius (r_m) – the radius from the core that contains half the total mass of the cluster. A small half-mass radius, relative to the overall size, indicates a dense core. Messier 3 (M3), for example, has an overall visible dimension of about 18 arc minutes, but a half-mass radius of only 1.12 arc minutes.

The tidal radius, or Hill sphere, is the distance from the center of the globular cluster at which the external gravitation of the galaxy has more influence over the stars in the cluster than does the cluster itself. This is the distance at which the individual stars belonging to a cluster can be separated away by the galaxy. The tidal radius of M3, for example, is about forty arc minutes, or about 113 pc.

Mass segregation, luminosity and core collapse

In most Milky Way clusters, the surface brightness of a globular cluster as a function of decreasing distance to the core first increases, then levels off at a distance typically 1–2 parsecs from the core. About 20% of the globular clusters have undergone a process termed "core collapse". The luminosity in such a cluster increases steadily all the way to the core region.

47 Tucanae is the second most luminous globular cluster in the Milky Way, after Omega Centauri.

Models of globular clusters predict that core collapse occurs when the more massive stars in a globular cluster encounter their less massive counterparts. Over time, dynamic processes cause individual stars to migrate from the center of the cluster to the outside, resulting in a net loss of kinetic energy from the core region and leading the region's remaining stars to occupy a more compact volume. When this gravothermal instability occurs, the central region of the cluster becomes densely crowded with stars, and the surface brightness of the cluster forms a power-law cusp. A massive black hole at the core could also result in a luminosity cusp. Over a long time, this leads to a concentration of massive stars near the core, a phenomenon called mass segregation.

The dynamical heating effect of binary star systems works to prevent an initial core collapse of the cluster. When a star passes near a binary system, the orbit of the latter pair tends to contract, releasing energy. Only after this primordial supply of energy is exhausted can a deeper core collapse proceed. In contrast, the effect of tidal shocks as a globular cluster repeatedly passes through the plane of a spiral galaxy tends to significantly accelerate core collapse.

Core collapse may be divided into three phases. During a cluster's adolescence, core collapse begins with stars nearest the core. Interactions between binary star systems prevents further collapse as the cluster approaches middle age. The central binaries are either disrupted or ejected, resulting in a tighter concentration at the core. The interaction of stars in the collapsed core region causes tight binary systems to form. As other stars interact with these tight binaries they increase the energy at the core, causing the cluster to re-expand. As the average time for a core collapse is typically less than the age of the galaxy, many of a galaxy's globular clusters may have passed through a core collapse stage, then re-expanded.

Globular cluster NGC 1854 is located in the Large Magellanic Cloud.

The HST has provided convincing observational evidence of this stellar mass-sorting process in globular clusters. Heavier stars slow down and crowd at the cluster's core, while lighter stars pick up speed and tend to spend more time at the cluster's periphery. The cluster 47 Tucanae, made up of about one million stars, is one of the densest globular clusters in the Southern Hemisphere. This cluster was subjected to an intensive photographic survey that obtained precise velocities for nearly fifteen thousand stars in this cluster.

The overall luminosities of the globular clusters within the Milky Way and the Andromeda Galaxy each have a roughly Gaussian distribution, with an average magnitude M_v and a variance σ². This distribution of globular cluster luminosities is called the Globular Cluster Luminosity Function (GCLF). For the Milky Way, M_v = −7.29 ± 0.13, σ = 1.1 ± 0.1. The GCLF has been used as a "standard candle" for measuring the distance to other galaxies, under the assumption that globular clusters in remote galaxies behave similarly to those in the Milky Way.

N-body simulations

Main article: N-body simulation

Computing the gravitational interactions between stars within a globular cluster requires solving the N-body problem. The naive computational cost for a dynamic simulation increases in proportion to N ² (where N is the number of objects), so the computing requirements to accurately simulate a cluster of thousands of stars can be enormous. A more efficient method of simulating the N-body dynamics of a globular cluster is done by subdivision into small volumes and velocity ranges, and using probabilities to describe the locations of the stars. Their motions are described by means of the Fokker–Planck equation, often using a model describing the mass density as a function of radius, such as a Plummer model. The simulation becomes more difficult when the effects of binaries and the interaction with external gravitation forces (such as from the Milky Way galaxy) must also be included. In 2010 a low-density globular cluster's lifetime evolution was able to be directly computed, star-by-star.

Completed N-body simulations have shown that stars can follow unusual paths through the cluster, often forming loops and falling more directly toward the core than would a single star orbiting a central mass. Additionally, some stars gain sufficient energy to escape the cluster due to gravitational interactions that result in a sufficient increase in velocity. Over long periods of time this process leads to the dissipation of the cluster, a process termed evaporation. The typical time scale for the evaporation of a globular cluster is 10¹⁰ years. The ultimate fate of a globular cluster must be either to accrete stars at its core, causing its steady contraction, or gradual shedding of stars from its outer layers.

Binary stars form a significant portion of stellar systems, with up to half of all field stars and open cluster stars occurring in binary systems. The present-day binary fraction in globular clusters is difficult to measure, and any information about their initial binary fraction is lost by subsequent dynamical evolution. Numerical simulations of globular clusters have demonstrated that binaries can hinder and even reverse the process of core collapse in globular clusters. When a star in a cluster has a gravitational encounter with a binary system, a possible result is that the binary becomes more tightly bound and kinetic energy is added to the solitary star. When the massive stars in the cluster are sped up by this process, it reduces the contraction at the core and limits core collapse.

Intermediate forms

Messier 10 lies about 15,000 light-years from Earth, in the constellation of Ophiuchus.

Cluster classification is not always definitive; objects have been found that can be classified in more than one category. For example, BH 176 in the southern part of the Milky Way has properties of both an open and a globular cluster.

In 2005 astronomers discovered a new, "extended" type of star cluster in the Andromeda Galaxy's halo, similar to the globular cluster. The three new-found clusters have a similar star count to globular clusters and share other characteristics, such as stellar populations and metallicity, but are distinguished by their larger size – several hundred light years across – and some hundred times lower density. Their stars are separated by larger distances; parametrically, these clusters lie somewhere between a globular cluster and a dwarf spheroidal galaxy. The formation of these extended clusters is likely related to accretion. It is unclear why the Milky Way lacks such clusters; Andromeda is unlikely to be the sole galaxy with them, but their presence in other galaxies remains unknown.

Tidal encounters

When a globular cluster comes close to a large mass, such as the core region of a galaxy, it undergoes a tidal interaction. The difference in gravitational strength between the nearer and further parts of the cluster results in an asymmetric, tidal force. A "tidal shock" occurs whenever the orbit of a cluster takes it through the plane of a galaxy.

Tidal shocks can pull stars away from the cluster halo, leaving only the core part of the cluster; these trails of stars can extend several degrees away from the cluster. These tails typically both precede and follow the cluster along its orbit and can accumulate significant portions of the original mass of the cluster, forming clump-like features. The globular cluster Palomar 5, for example, is near the apogalactic point of its orbit after passing through the Milky Way. Streams of stars extend outward toward the front and rear of the orbital path of this cluster, stretching to distances of 13,000 light years. Tidal interactions have stripped away much of Palomar 5's mass; further interactions with the galactic core are expected to transform it into a long stream of stars orbiting the Milky Way in its halo.

The Milky Way is in the process of tidally stripping the Sagittarius Dwarf Spheroidal Galaxy of stars and globular clusters through the Sagittarius Stream. As many as 20% of the globular clusters in the Milky Way's outer halo may have originated in that galaxy. Palomar 12, for example, most likely originated in the Sagittarius Dwarf Spheroidal but is now associated with the Milky Way. Tidal interactions like these add kinetic energy into a globular cluster, dramatically increasing the evaporation rate and shrinking the size of the cluster. The increased evaporation accelerates the process of core collapse.

Planets

Astronomers are searching for exoplanets of stars in globular star clusters. A search in 2000 for giant planets in the globular cluster 47 Tucanae came up negative, suggesting that the abundance of heavier elements – low in globular clusters – necessary to build these planets may need to be at least 40% of the Sun's abundance. Because terrestrial planets are built from heavier elements such as silicon, iron and magnesium, member stars have a far lower likelihood of hosting Earth-mass planets than stars in the solar neighborhood. Globular clusters are thus unlikely to host habitable terrestrial planets.

A giant planet was found in the globular cluster Messier 4, orbiting a pulsar in the binary star system PSR B1620-26. The planet's eccentric and highly inclined orbit suggests it may have been formed around another star in the cluster, then "exchanged" into its current arrangement. The likelihood of close encounters between stars in a globular cluster can disrupt planetary systems; some planets break free to become rogue planets, orbiting the galaxy. Planets orbiting close to their star can become disrupted, potentially leading to orbital decay and an increase in orbital eccentricity and tidal effects. In 2024, a gas giant or brown dwarf was found to closely orbit the pulsar "M62H", where the name indicates that the planetary system belongs to the globular cluster Messier 62.