Type II supernova

From Wikipedia, the free encyclopedia

 

The expanding remnant of SN 1987A, a Type II-P supernova in the Large Magellanic Cloud. NASA image.

A Type II supernova (plural: supernovae) results from the rapid collapse and violent explosion of a massive star. A star must have at least 8 times, and no more than 40–50 times, the mass of the Sun for this type of explosion.^[1] It is distinguished from other types of supernovae by the presence of hydrogen in its spectrum. Type II supernovae are mainly observed in the spiral arms of galaxies and in H II regions, but not in elliptical galaxies.

Stars generate energy by the nuclear fusion of elements. Unlike the Sun, massive stars possess the mass needed to fuse elements that have an atomic mass greater than hydrogen and helium, albeit at increasingly higher temperatures and pressures, causing increasingly shorter stellar life spans. The degeneracy pressure of electrons and the energy generated by these fusion reactions are sufficient to counter the force of gravity and prevent the star from collapsing, maintaining stellar equilibrium. The star fuses increasingly higher mass elements, starting with hydrogen and then helium, progressing up through the periodic table until a core of iron and nickel is produced. Fusion of iron or nickel produces no net energy output, so no further fusion can take place, leaving the nickel-iron core inert. Due to the lack of energy output allowing outward pressure, equilibrium is broken.

When the mass of the inert core exceeds the Chandrasekhar limit of about 1.4 solar masses, electron degeneracy alone is no longer sufficient to counter gravity and maintain stellar equilibrium. A cataclysmic implosion takes place within seconds, in which the outer core reaches an inward velocity of up to 23% of the speed of light and the inner core reaches temperatures of up to 100 billion kelvin. Neutrons and neutrinos are formed via reversed beta-decay, releasing about 10⁴⁶ joules (100 foes) in a ten-second burst. The collapse is halted by neutron degeneracy, causing the implosion to rebound and bounce outward. The energy of this expanding shock wave is sufficient to accelerate the surrounding stellar material to escape velocity, forming a supernova explosion, while the shock wave and extremely high temperature and pressure briefly allow for the production of elements heavier than iron.^[2] Depending on initial size of the star, the remnants of the core form a neutron star or a black hole. Because of the underlying mechanism, the resulting nova is also described as a core-collapse supernova.

There exist several categories of Type II supernova explosions, which are categorized based on the resulting light curve—a graph of luminosity versus time—following the explosion. Type II-L supernovae show a steady (linear) decline of the light curve following the explosion, whereas Type II-P display a period of slower decline (a plateau) in their light curve followed by a normal decay. Type Ib and Ic supernovae are a type of core-collapse supernova for a massive star that has shed its outer envelope of hydrogen and (for Type Ic) helium. As a result, they appear to be lacking in these elements.

Formation

The onion-like layers of a massive, evolved star just before core collapse. (Not to scale.)

Stars far more massive than the sun evolve in more complex ways. In the core of the star, hydrogen is fused into helium, releasing thermal energy that heats the sun's core and provides outward pressure that supports the sun's layers against collapse in a process known as stellar or hydrostatic equilibrium. The helium produced in the core accumulates there since temperatures in the core are not yet high enough to cause it to fuse. Eventually, as the hydrogen at the core is exhausted, fusion starts to slow down, and gravity causes the core to contract. This contraction raises the temperature high enough to initiate a shorter phase of helium fusion, which accounts for less than 10% of the star's total lifetime. In stars with fewer than eight solar masses, the carbon produced by helium fusion does not fuse, and the star gradually cools to become a white dwarf.^[3][4] White dwarf stars, if they have a near companion, may then become Type Ia supernovae.

A much larger star, however, is massive enough to create temperatures and pressures needed to cause the carbon in the core to begin to fuse once the star contracts at the end of the helium-burning stage. The cores of these massive stars become layered like onions as progressively heavier atomic nuclei build up at the center, with an outermost layer of hydrogen gas, surrounding a layer of hydrogen fusing into helium, surrounding a layer of helium fusing into carbon via the triple-alpha process, surrounding layers that fuse to progressively heavier elements. As a star this massive evolves, it undergoes repeated stages where fusion in the core stops, and the core collapses until the pressure and temperature are sufficient to begin the next stage of fusion, reigniting to halt collapse.^[3][4]

Core-burning nuclear fusion stages for a 25-solar mass star

Process	Main fuel	Main products	25 M☉ star[5]
			Temperature (Kelvin)	Density (g/cm3)	Duration
hydrogen burning	hydrogen	helium	7×107	10	107 years
triple-alpha process	helium	carbon, oxygen	2×108	2000	106 years
carbon burning process	carbon	Ne, Na, Mg, Al	8×108	106	103 years
neon burning process	neon	O, Mg	1.6×109	107	3 years
oxygen burning process	oxygen	Si, S, Ar, Ca	1.8×109	107	0.3 years
silicon burning process	silicon	nickel (decays into iron)	2.5×109	108	5 days

Core collapse

The factor limiting this process is the amount of energy that is released through fusion, which is dependent on the binding energy that holds together these atomic nuclei. Each additional step produces progressively heavier nuclei, which release progressively less energy when fusing. In addition, from carbon-burning onwards, energy loss via neutrino production becomes significant, leading to a higher rate of reaction than would otherwise take place.^[6] This continues until nickel-56 is produced, which decays radioactively into cobalt-56 and then iron-56 over the course of a few months. As iron and nickel have the highest binding energy per nucleon of all the elements,^[7] energy cannot be produced at the core by fusion, and a nickel-iron core grows.^[4][8] This core is under huge gravitational pressure. As there is no fusion to further raise the star's temperature to support it against collapse, it is supported only by degeneracy pressure of electrons. In this state, matter is so dense that further compaction would require electrons to occupy the same energy states. However, this is forbidden for identical fermion particles, such as the electron – a phenomenon called the Pauli exclusion principle.

When the core's mass exceeds the Chandrasekhar limit of about 1.4 solar masses, degeneracy pressure can no longer support it, and catastrophic collapse ensues.^[9] The outer part of the core reaches velocities of up to 70,000 km/s (23% of the speed of light) as it collapses toward the center of the star.^[10] The rapidly shrinking core heats up, producing high-energy gamma rays that decompose iron nuclei into helium nuclei and free neutrons via photodisintegration. As the core's density increases, it becomes energetically favorable for electrons and protons to merge via inverse beta decay, producing neutrons and elementary particles called neutrinos. Because neutrinos rarely interact with normal matter, they can escape from the core, carrying away energy and further accelerating the collapse, which proceeds over a timescale of milliseconds. As the core detaches from the outer layers of the star, some of these neutrinos are absorbed by the star's outer layers, beginning the supernova explosion.^[11]

For Type II supernovae, the collapse is eventually halted by short-range repulsive neutron-neutron interactions, mediated by the strong force, as well as by degeneracy pressure of neutrons, at a density comparable to that of an atomic nucleus. Once collapse stops, the infalling matter rebounds, producing a shock wave that propagates outward. The energy from this shock dissociates heavy elements within the core. This reduces the energy of the shock, which can stall the explosion within the outer core.^[12]

The core collapse phase is so dense and energetic that only neutrinos are able to escape. As the protons and electrons combine to form neutrons by means of electron capture, an electron neutrino is produced. In a typical Type II supernova, the newly formed neutron core has an initial temperature of about 100 billion kelvin, 10⁴ times the temperature of the sun's core. Much of this thermal energy must be shed for a stable neutron star to form, otherwise the neutrons would "boil away". This is accomplished by a further release of neutrinos.^[13] These 'thermal' neutrinos form as neutrino-antineutrino pairs of all flavors, and total several times the number of electron-capture neutrinos.^[14] The two neutrino production mechanisms convert the gravitational potential energy of the collapse into a ten second neutrino burst, releasing about 10⁴⁶ joules (100 foes).^[15]

Through a process that is not clearly understood, about 10⁴⁴ joules (1 foe) is reabsorbed by the stalled shock, producing an explosion.^[a][12] The neutrinos generated by a supernova were actually observed in the case of Supernova 1987A, leading astronomers to conclude that the core collapse picture is basically correct. The water-based Kamiokande II and IMB instruments detected antineutrinos of thermal origin,^[13] while the gallium-71-based Baksan instrument detected neutrinos (lepton number = 1) of either thermal or electron-capture origin.

Within a massive, evolved star (a) the onion-layered shells of elements undergo fusion, forming a nickel-iron core (b) that reaches Chandrasekhar-mass and starts to collapse. The inner part of the core is compressed into neutrons (c), causing infalling material to bounce (d) and form an outward-propagating shock front (red). The shock starts to stall (e), but it is re-invigorated by neutrino interaction. The surrounding material is blasted away (f), leaving only a degenerate remnant.

When the progenitor star is below about 20 solar masses – depending on the strength of the explosion and the amount of material that falls back – the degenerate remnant of a core collapse is a neutron star.^[10] Above this mass, the remnant collapses to form a black hole.^[4][16] The theoretical limiting mass for this type of core collapse scenario is about 40–50 solar masses. Above that mass, a star is believed to collapse directly into a black hole without forming a supernova explosion,^[17] although uncertainties in models of supernova collapse make calculation of these limits uncertain.

Theoretical models

The Standard Model of particle physics is a theory which describes three of the four known fundamental interactions between the elementary particles that make up all matter. This theory allows predictions to be made about how particles will interact under many conditions. The energy per particle in a supernova is typically one to one hundred and fifty picojoules (tens to hundreds of MeV).^[18] The per-particle energy involved in a supernova is small enough that the predictions gained from the Standard Model of particle physics are likely to be basically correct. But the high densities may require corrections to the Standard Model.^[19] In particular, Earth-based particle accelerators can produce particle interactions which are of much higher energy than are found in supernovae,^[20] but these experiments involve individual particles interacting with individual particles, and it is likely
that the high densities within the supernova will produce novel effects. The interactions between neutrinos and the other particles in the supernova take place with the weak nuclear force, which is believed to be well understood. However, the interactions between the protons and neutrons involve the strong nuclear force, which is much less well understood.^[21]

The major unsolved problem with Type II supernovae is that it is not understood how the burst of neutrinos transfers its energy to the rest of the star producing the shock wave which causes the star to explode. From the above discussion, only one percent of the energy needs to be transferred to produce an explosion, but explaining how that one percent of transfer occurs has proven very difficult, even though the particle interactions involved are believed to be well understood. In the 1990s, one model for doing this involved convective overturn, which suggests that convection, either from neutrinos from below, or infalling matter from above, completes the process of destroying the progenitor star. Heavier elements than iron are formed during this explosion by neutron capture, and from the pressure of the neutrinos pressing into the boundary of the "neutrinosphere", seeding the surrounding space with a cloud of gas and dust which is richer in heavy elements than the material from which the star originally formed.^[22]

Neutrino physics, which is modeled by the Standard Model, is crucial to the understanding of this process.^[19] The other crucial area of investigation is the hydrodynamics of the plasma that makes up the dying star; how it behaves during the core collapse determines when and how the "shock wave" forms and when and how it "stalls" and is reenergized.^[23]

In fact, some theoretical models incorporate a hydrodynamical instability in the stalled shock known as the "Standing Accretion Shock Instability" (SASI). This instability comes about as a consequence of non-spherical perturbations oscillating the stalled shock thereby deforming it. The SASI is often used in tandem with neutrino theories in computer simulations for re-energizing the stalled shock.^[24]

Computer models have been very successful at calculating the behavior of Type II supernovae once the shock has been formed. By ignoring the first second of the explosion, and assuming that an explosion is started, astrophysicists have been able to make detailed predictions about the elements produced by the supernova and of the expected light curve from the supernova.^[25][26][27]

Light curves for Type II-L and Type II-P supernovae

This graph of the luminosity as a function of time shows the characteristic shapes of the light curves for a Type II-L and II-P supernova.

When the spectrum of a Type II supernova is examined, it normally displays Balmer absorption lines – reduced flux at the characteristic frequencies where hydrogen atoms absorb energy. The presence of these lines is used to distinguish this category of supernova from a Type I supernova.

When the luminosity of a Type II supernova is plotted over a period of time, it shows a characteristic rise to a peak brightness followed by a decline. These light curves have an average decay rate of 0.008 magnitudes per day; much lower than the decay rate for Type Ia supernovae. Type II are sub-divided into two classes, depending on the shape of the light curve. The light curve for a Type II-L supernova shows a steady (linear) decline following the peak brightness. By contrast, the light curve of a Type II-P supernova has a distinctive flat stretch (called a plateau) during the decline; representing a period where the luminosity decays at a slower rate. The net luminosity decay rate is lower, at 0.0075 magnitudes per day for Type II-P, compared to 0.012 magnitudes per day for Type II-L.^[28]

The difference in the shape of the light curves is believed to be caused, in the case of Type II-L supernovae, by the expulsion of most of the hydrogen envelope of the progenitor star.^[28] The plateau phase in Type II-P supernovae is due to a change in the opacity of the exterior layer. The shock wave ionizes the hydrogen in the outer envelope – stripping the electron from the hydrogen atom – resulting in a significant increase in the opacity. This prevents photons from the inner parts of the explosion from escaping. Once the hydrogen cools sufficiently to recombine, the outer layer becomes transparent.^[29]

Type IIn supernovae

The "n" denotes narrow, which indicates the presence of intermediate or very narrow width H emission lines in the spectra. In the intermediate width case, the ejecta from the explosion may be interacting strongly with gas around the star – the circumstellar medium. ^[30][31] There are indications that they originate as stars similar to Luminous blue variables with large mass losses before exploding.^[32] SN 2005gl is one example of Type IIn; SN 2006gy, an extremely energetic supernova, may be another example.^[33]

Type IIb supernovae

A Type IIb supernova has a weak hydrogen line in its initial spectrum, which is why it is classified as a Type II. However, later on the H emission becomes undetectable, and there is also a second peak in the light curve that has a spectrum which more closely resembles a Type Ib supernova. The progenitor could have been a giant star which lost most of its hydrogen envelope due to interactions with a companion in a binary system, leaving behind the core that consisted almost entirely of helium.^[34] As the ejecta of a Type IIb expands, the hydrogen layer quickly becomes more transparent and reveals the deeper layers.^[34] The classic example of a Type IIb supernova is Supernova 1993J,^[35][36] while another example is Cassiopeia A.^[37] The IIb class was first introduced (as a theoretical concept) by Ensman & Woosley 1987.

Hypernovae (collapsars)

Hypernovae are a rare type of supernova substantially more luminous and energetic than standard supernovae. Examples are 1997ef (type Ic) and 1997cy (type IIn). Hypernovae are produced by more than one type of event: relativistic jets during formation of a black hole from fallback of material onto the neutron star core, the collapsar model; interaction with a dense envelope of circumstellar material, the CSM model; the highest mass pair instability supernovae; possibly others such as binary and quark star model.
Stars with initial masses between about 25 and 90 times the sun develop cores large enough that after a supernova explosion, some material will fall back onto the neutron star core and create a black hole. In many cases this reduces the luminosity of the supernova, and above 90 masses the star collapses directly into a black hole without a supernova explosion. However if the progenitor is spinning quickly enough the infalling material generates relativistic jets that emit more energy than the original explosion.^[38] They may also be seen directly if beamed towards us, giving the impression of an even more luminous object. In some cases these can produce gamma-ray bursts, although not all gamma-ray bursts are from supernovae.^[39]

In some cases a type II supernova occurs when the star is surrounded by a very dense cloud of material, most likely expelled during luminous blue variable eruptions. This material is shocked by the explosion and becomes more luminous than a standard supernova. It is likely that there is a range of luminosities for these type IIn supernovae with only the brightest qualifying as a hypernova.

Pair instability supernovae occur when an oxygen core in an extremely massive star becomes hot enough that gamma rays spontaneously produce electron-positron pairs.^[40] This causes the core to collapse, but where the collapse of an iron core causes endothermic fusion to heavier elements, the collapse of an oxygen core creates runaway exothermic fusion which completely unbinds the star. The total energy emitted depends on the initial mass, with much of the core being converted to ⁵⁶Ni and ejected which then powers the supernova for many months. At the lower end stars of about 140 solar masses produce supernovae that are long-lived but otherwise typical, while the highest mass stars of around 250 solar masses produce supernovae that are extremely luminous and also very long lived; hypernovae. More massive stars die by photodisintegration. Only population III stars, with very low metallicity, can reach this stage. Stars with more heavy elements are more opaque and blow away their outer layers until they are small enough to explode as a normal type Ib/c supernova. It is thought that even in our own galaxy, mergers of old low metallicity stars may form massive stars capable of creating a pair instability supernova.

Planetary nebula

From Wikipedia, the free encyclopedia

 

X-ray/optical composite image of the Cat's Eye Nebula.

NGC 6326, a planetary nebula with glowing wisps of outpouring gas that are lit up by a binary^[1] central star.

A planetary nebula, often abbreviated as PN or plural PNe, is a kind of emission nebula consisting of an expanding glowing shell of ionized gas ejected from old red giant stars late in their lives.^[2] The word 'nebula' is Latin for mist or cloud and the term 'planetary nebula' is a misnomer that originated in the 1780s with astronomer William Herschel because when viewed through his telescope, these objects appeared to him to resemble the rounded shapes of planets. Herschel's name for these objects was popularly adopted and has not been changed.^[3][4] They are a relatively short-lived phenomenon, lasting a few tens of thousands of years, compared to a typical stellar lifetime of several billion years.

A mechanism for formation of most planetary nebulae is thought to be the following: at the end of the star's life, during the red giant phase, the outer layers of the star are expelled by strong stellar winds. Eventually, after most of the red giant's atmosphere is dissipated, the exposed hot, luminous core emits ultraviolet radiation to ionize the ejected outer layers of the star.^[2] Absorbed ultraviolet light energises the shell of nebulous gas around the central star, appearing as a bright coloured planetary nebula at several discrete visible wavelengths.

Planetary nebulae may play a crucial role in the chemical evolution of the Milky Way, returning material to the interstellar medium from stars where elements, the products of nucleosynthesis (such as carbon, nitrogen, oxygen and neon), have been created. Planetary nebulae are also observed in more distant galaxies, yielding useful information about their chemical abundances.

In recent years, Hubble Space Telescope images have revealed many planetary nebulae to have extremely complex and varied morphologies. About one-fifth are roughly spherical, but the majority are not spherically symmetric. The mechanisms which produce such a wide variety of shapes and features are not yet well understood, but binary central stars, stellar winds and magnetic fields may play a role.

Observations

NGC 7293, The Helix Nebula
Credit: NASA, ESA, and C.R. O'Dell (Vanderbilt University)

NGC 2392, The Eskimo Nebula
Credit: NASA, ESA, Andrew Fruchter (STScI), and the ERO team (STScI + ST-ECF)

Planetary nebulae are generally faint objects; none are visible to the naked eye. The first planetary nebula discovered was the Dumbbell Nebula in the constellation of Vulpecula. It was observed by Charles Messier in 1764 and listed as M27 in his catalogue of nebulous objects.^[5] To early observers with low-resolution telescopes, M27 and subsequently discovered planetary nebulae somewhat resembled the giant planets like Uranus. William Herschel, discoverer of Uranus, eventually coined the term 'planetary nebula' for them.^[5][6] At first Herschel thought the objects were stars surrounded by material that was condensing into planets rather than what is known to be evidence of dead stars that have incinerated any orbiting planets.^[7]

The nature of planetary nebulae was unknown until the first spectroscopic observations were made in the mid-19th century. Using a prism to disperse their light, William Huggins was one of the earliest astronomers to study the optical spectra of astronomical objects.^[6] On August 29, 1864, Huggins was the first to analyze the spectrum of a planetary nebula when he observed NGC 6543.^[5] His observations of stars showed that their spectra consisted of a continuum of radiation with many dark lines superimposed on them. He later found that many nebulous objects such as the Andromeda Nebula (as it was then known) had spectra that were quite similar. These nebulae were later shown to be galaxies.

However, when Huggins looked at the Cat's Eye Nebula, he found a very different spectrum. Rather than a strong continuum with absorption lines superimposed, the Cat's Eye Nebula and other similar objects showed only a small number of emission lines.^[6] The brightest of these was at a wavelength of 500.7 nanometres, which did not correspond with a line of any known element.^[8] At first it was hypothesized that the line might be due to an unknown element, which was named nebulium. A similar idea had led to the discovery of helium through analysis of the Sun's spectrum in 1868.^[5]

While helium was isolated on earth soon after its discovery in the spectrum of the sun, nebulium was not. In the early 20th century Henry Norris Russell proposed that rather than being a new element, the line at 500.7 nm was due to a familiar element in unfamiliar conditions.^[5]

Physicists showed in the 1920s that in gas at extremely low densities, electrons can populate excited metastable energy levels in atoms and ions which at higher densities are rapidly de-excited by collisions.^[9] Electron transitions from these levels in nitrogen and oxygen ions (O⁺, O²⁺ or OIII, and N⁺) give rise to the 500.7 nm emission line and others.^[5] These spectral lines, which can only be seen in very low density gases, are called forbidden lines. Spectroscopic observations thus showed that nebulae were made of extremely rarefied gas.^[10]

The central stars of planetary nebulae are very hot.^[2] Only once a star has exhausted most of its nuclear fuel can it collapse to such a small size. Planetary nebulae came to be understood as a final stage of stellar evolution. Spectroscopic observations show that all planetary nebulae are expanding. This led to the idea that planetary nebulae were caused by a star's outer layers being thrown into space at the end of its life.^[5]

Towards the end of the 20th century, technological improvements helped to further the study of planetary nebulae.^[11] Space telescopes allowed astronomers to study light wavelengths outside those that the earth's atmosphere transmits. Infrared and ultraviolet studies of planetary nebulae allowed much more accurate determinations of nebular temperatures, densities and elemental abundances.^[12][13] Charge-coupled device technology allowed much fainter spectral lines to be measured accurately than had previously been possible. The Hubble Space Telescope also showed that while many nebulae appear to have simple and regular structures when observed from the ground, the very high optical resolution achievable by telescopes above the Earth's atmosphere reveals extremely complex structures.^[14][15]

Under the Morgan-Keenan spectral classification scheme, planetary nebulae are classified as Type-P, although this notation is seldom used in practice.^[16]

Origins

Computer simulation of the formation of a planetary nebula from a star with a warped disk, showing the complexity which can result from a small initial asymmetry.
Credit: Vincent Icke

Stars greater than 8 solar masses (M_⊙) will likely end their lives in dramatic supernovae explosions, while planetary nebulae seemingly only occur at the end of the lives of intermediate and low mass stars between 0.8 M_⊙ to 8.0 M_⊙. ^[17] Progenitor stars that form planetary nebulae will spend most of their lifetimes converting their hydrogen into helium in the star's core by nuclear fusion at about 15 million K. This generated energy creates outward pressure from fusion reactions in the core, equally balancing the crushing inward pressures of the star's gravity.^[18] Hence, all single intermediate to low-mass stars on the main sequence can last for tens of millions to billions of years.

When the hydrogen source in the core starts to diminish, gravity starts compressing the core, causing a rise in temperature to about 100 million K.^[19] Such higher core temperatures then make the star's cooler outer layers expand to create much larger red giant stars. This end phase causes a dramatic rise in stellar luminosity, where the released energy is distributed over a much larger surface area, even though the average surface temperature is lower. In stellar evolution terms, stars undergoing such increases in luminosity are known as asymptotic giant branch stars (AGB).^[19]

For the more massive asymptotic giant branch stars that form planetary nebulae, whose progenitors exceed about 3M_⊙, their cores will continue to contract. When temperatures reach about 100 million K, the available helium nuclei fuse into carbon and oxygen, so that the star again resumes to radiate energy, temporarily stopping the core's contraction. This new helium burning phase (fusion of helium nuclei) forms a growing inner core of inert carbon and oxygen. Above it is a thin helium-burning shell, surrounded in turn by a hydrogen-burning shell. However, this new phase lasts only 20,000 years or so, a short period compared to the entire lifetime of the star.

In either scenario, the venting of atmosphere continues unabated into interstellar space, but when the outer surface of the exposed core reaches temperatures exceeding about 30,000 K, there are enough emitted ultraviolet photons to ionize the ejected atmosphere, causing the gas to shine as a planetary nebula.^[19]

Lifetime

The Necklace Nebula consists of a bright ring, measuring about two light-years across, dotted with dense, bright knots of gas that resemble diamonds in a necklace. The knots glow brightly due to absorption of ultraviolet light from the central stars.^[20]

After a star passes through the asymptotic giant branch (AGB) phase, the short planetary nebula phase of stellar evolution begins^[11] as gases blown away from the central star at speeds of a few kilometers per second. The central star is the remnant of its AGB progenitor, an electron-degenerate carbon-oxygen core that has lost most of its hydrogen envelope due to mass loss on the AGB.^[11] As the gases expand, the central star undergoes a two-stage evolution, first growing hotter as it continues to contract and hydrogen fusion reactions occur in the shell around the core and then slowly cooling once the hydrogen shell is exhausted through fusion and mass loss.^[11] In the second phase, it radiates away its energy and fusion reactions cease, as the central star is not heavy enough to generate the core temperatures required for carbon and oxygen to fuse.^[5][11] During the first phase, the central star maintains constant luminosity,^[11] while at the same time it grows ever hotter, eventually reaching temperatures around 100,000 K. In the second phase, it cools so much that it does not give off enough ultraviolet radiation to ionize the increasingly distant gas cloud. The star becomes a white dwarf, and the expanding gas cloud becomes invisible to us, ending the planetary nebula phase of evolution.^[11] For a typical planetary nebula, about 10,000 years^[11] passes between its formation and recombination of the star.^[5]

Galactic recyclers

Planetary nebulae play a very important role in galactic evolution. The early universe consisted almost entirely of hydrogen and helium, but stars create heavier elements via nuclear fusion. The gases of planetary nebulae thus contain a large proportion of elements such as carbon, nitrogen and oxygen, and as they expand and merge into the interstellar medium, they enrich it with these heavy elements, collectively known as metals by astronomers.^[21]

Subsequent generations of stars which form from the nebular remains will then have a higher initial content of heavier elements. Even though the heavy elements will still be a very small component of the star, they have a marked effect on its evolution. Stars which formed very early in the universe and contain small quantities of heavy elements are known as Population II stars, while younger stars with higher heavy element content are known as Population I stars (see stellar population).^[22]

Characteristics

Physical characteristics

NGC 6720, The Ring Nebula
Credit: STScI/AURA

Lemon slice nebula (IC 3568).

A typical planetary nebula is roughly one light year across, and consists of extremely rarefied gas, with a density generally from 100 to 10,000 particles per cm³.^[23] (The Earth's atmosphere, by comparison, contains 2.5×10¹⁹ particles per cm³.) Young planetary nebulae have the highest densities, sometimes as high as 10⁶ particles per cm³. As nebulae age, their expansion causes their density to decrease. The masses of planetary nebulae range from 0.1 to 1 solar masses.^[23]

Radiation from the central star heats the gases to temperatures of about 10,000 K.^[24] The gas temperature in central regions is usually much higher than at the periphery reaching 16,000–25,000 K.^[25] The volume in the vicinity of the central star is often filled with a very hot (coronal) gas having the temperature of about 1,000,000 K. This gas originates from the surface of the central star in the form of the fast stellar wind.^[26]

Nebulae may be described as matter bounded or radiation bounded. In the former case, there is not enough matter in the nebula to absorb all the UV photons emitted by the star, and the visible nebula is fully ionized. In the latter case, there are not enough UV photons being emitted by the central star to ionize all the surrounding gas, and an ionization front propagates outward into the circumstellar envelope of neutral atoms.^[27]

Numbers and distribution

About 3000 planetary nebulae are now known to exist in our galaxy,^[28] out of 200 billion stars. Their very short lifetime compared to total stellar lifetime accounts for their rarity. They are found mostly near the plane of the Milky Way, with the greatest concentration near the galactic center.^[29]

Morphology

Play media

This animation shows how the two stars at the heart of a planetary nebula like Fleming 1 can control the creation of the spectacular jets of material ejected from the object.

Only about 20% of planetary nebulae are spherically symmetric (for example, see Abell 39).^[30] A wide variety of shapes exist with some very complex forms seen. Planetary nebulae are classified by different authors into: stellar, disk, ring, irregular, helical, bipolar, quadrupolar,^[31] and other types,^[32] although the majority of them belong to just three types: spherical, elliptical and bipolar. Bipolar nebulae are concentrated in the galactic plane, likely produced by relatively young massive progenitor stars; and bipolars in the galactic bulge appear to prefer orienting their orbital axes parallel to the galactic plane.^[33] On the other hand, spherical nebulae are likely produced by the old stars similar to the Sun.^[26]

The huge variety of the shapes is partially the projection effect—the same nebula when viewed under different angles will appear different. Nevertheless, the reason for the huge variety of physical shapes is not fully understood.^[32] Gravitational interactions with companion stars if the central stars are double stars may be one cause. Another possibility is that planets disrupt the flow of material away from the star as the nebula forms. It has been determined that the more massive stars produce more irregularly shaped nebulae.^[34] In January 2005, astronomers announced the first detection of magnetic fields around the central stars of two planetary nebulae, and hypothesized that the fields might be partly or wholly responsible for their remarkable shapes.^[35][36]

Membership in clusters

Abell 78, 24 inch telescope on Mt. Lemmon, AZ. Courtesy of Joseph D. Schulman.

Planetary nebulae have been detected as members in four globular clusters: Messier 15, Messier 22, NGC 6441 and Palomar 6. However, there is currently only one established case of a planetary nebula discovered in an open cluster.^[37] The cases of NGC 2348 in Messier 46, and NGC 2818, are often cited as bona fide instances, however, they are instead line-of-sight coincidences.^[29][38][39]

Partly because of their small total mass, open clusters have relatively poor gravitational cohesion. Consequently, open clusters tend to disperse after a relatively short time, typically from 100 to 600 million years.^[40]

Theoretical models predict that planetary nebulae can form from main-sequence stars of between one and eight solar masses, which puts the progenitor star's age at greater than 40 million years. Although there are a few hundred known open clusters within that age range, a variety of reasons limit the chances of finding a planetary nebula within. For one reason, the planetary nebula phase for more massive stars is on the order of thousands of years—a blink of the eye in cosmic terms.^[29]

Current issues in planetary nebula studies

Odd pair of aging stars sculpt spectacular shape of planetary nebula.^[41]

Tiny planetary nebula NGC 6886.

The distances to planetary nebulae are generally poorly determined.^[42] It is possible to determine distances to the nearest planetary nebula by measuring their expansion rates. High resolution observations taken several years apart will show the expansion of the nebula perpendicular to the line of sight, while spectroscopic observations of the Doppler shift will reveal the velocity of expansion in the line of sight. Comparing the angular expansion with the derived velocity of expansion will reveal the distance to the nebula.^[14]

The issue of how such a diverse range of nebular shapes can be produced is a debatable topic. It is theorised that interactions between material moving away from the star at different speeds gives rise to most observed shapes.^[32] However, some astronomers postulate that close binary central stars might be responsible for the more complex and extreme planetary nebulae.^[43] Several have been shown to exhibit strong magnetic fields,^[44] and their interactions with ionized gas could explain some planetary nebulae shapes.^[36]

There are two main methods of determining metal abundances in nebulae. These rely on recombination lines and collisionally excited lines. Large discrepancies are sometimes seen between the results derived from the two methods. This may be explained by the presence of small temperature fluctuations within planetary nebulae. The discrepancies may be too large to be caused by temperature effects, and some hypothesize the existence of cold knots containing very little hydrogen to explain the observations. However, such knots have yet to be observed.^[45]

Globular cluster

From Wikipedia, the free encyclopedia

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.^[1]

A globular cluster is a spherical collection of stars that orbits a galactic core as a satellite. Globular clusters are very tightly bound by gravity, which gives them their spherical shapes and relatively high stellar densities toward their centers. The name of this category of star cluster is derived from the Latin globulus—a small sphere. A globular cluster is sometimes known more simply as a globular.

Globular clusters, which are found in the halo of a galaxy, contain considerably more stars and are much older than the less dense galactic, or open clusters, which are found in the disk. Globular clusters are fairly common; there are about 150^[2] to 158^[3] currently known globular clusters in the Milky Way, with perhaps 10 to 20 more still undiscovered.^[4] Large galaxies can have more: Andromeda, for instance, may have as many as 500.^[5] Some giant elliptical galaxies, particularly those at the centers of galaxy clusters, such as M87,^[6] have as many as 13,000 globular clusters.

These globular clusters orbit the galaxy out to large radii, 40 kiloparsecs (approximately 131,000 light-years) or more.^[7]

Every galaxy of sufficient mass in the Local Group has an associated group of globular clusters, and almost every large galaxy surveyed has been found to possess a system of globular clusters.^[8] The Sagittarius Dwarf galaxy and the disputed Canis Major Dwarf galaxy appear to be in the process of donating their associated globular clusters (such as Palomar 12) to the Milky Way.^[9] This demonstrates how many of this galaxy's globular clusters might have been acquired in the past.

Although it appears that globular clusters contain some of the first stars to be produced in the galaxy, their origins and their role in galactic evolution are still unclear. It does appear clear that globular clusters are significantly different from dwarf elliptical galaxies and were formed as part of the star formation of the parent galaxy rather than as a separate galaxy.^[10] However, recent conjectures by astronomers suggest that globular clusters and dwarf spheroidals may not be clearly separate and distinct types of objects.^[11]

Observation history

The first globular cluster discovered was M22 in 1665 by Abraham Ihle, a German amateur astronomer.^[12] However, given the small aperture of early telescopes, individual stars within a globular cluster were not resolved until Charles Messier observed M4.^[13] The first eight globular clusters discovered are shown in the table. Subsequently, Abbé Lacaille would list NGC 104, NGC 4833, M55, M69, and NGC 6397 in his 1751–52 catalogue. The M before a number refers to the catalogue of Charles Messier, while NGC is from the New General Catalogue by John Dreyer.
William Herschel began a survey program in 1782 using larger telescopes and was able to resolve the stars in all 33 of the known globular clusters. In addition he found 37 additional clusters. In Herschel's 1789 catalog of deep sky objects, his second such, he became the first to use the name globular cluster as their description.^[13]

The number of globular clusters discovered continued to increase, reaching 83 in 1915, 93 in 1930 and 97 by 1947. A total of 152 globular clusters have now been discovered in the Milky Way galaxy, out of an estimated total of 180 ± 20.^[4] These additional, undiscovered globular clusters are believed to be hidden behind the gas and dust of the Milky Way.

Beginning in 1914, Harlow Shapley began a series of studies of globular clusters, published in about 40 scientific papers. He examined the RR Lyrae variables in the clusters (which he assumed were cepheid variables) and would use their period–luminosity relationship for distance estimates. Later, it was found that RR Lyrae variables are fainter than cepheid variables, which caused Shapley to overestimate the distance to the clusters.^[14]

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

Of the globular clusters within our Milky Way, the majority are found in the vicinity of the galactic core, and the large majority lie on the side of the celestial sky centered on the core. In 1918, this strongly asymmetrical distribution was used by Harlow Shapley to make a determination of the overall dimensions of the galaxy. By 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.^[15] While his distance estimate was significantly in error, it did demonstrate that the dimensions of the galaxy were much greater than had been previously thought. His error was because dust in the Milky Way diminished the amount of light from a globular cluster that reached the earth, thus making it appear farther away. Shapley's estimate was, however, within the same order of magnitude as the currently accepted value.

Shapley's measurements also indicated that the Sun was relatively far from the center of the galaxy, contrary to what had previously been inferred from the apparently nearly even distribution of ordinary stars. In reality, ordinary stars lie within the galaxy's disk and are thus often obscured by gas and dust, whereas globular clusters lie outside the disk and can be seen at much further distances.

Classification of globulars

Shapley was subsequently assisted in his studies of clusters by Henrietta Swope and Helen Battles Sawyer (later Hogg). In 1927–29, Harlow Shapley and Helen Sawyer began categorizing clusters according to the degree of concentration the system has toward the core. The most concentrated clusters were identified as Class I, with successively diminishing concentrations ranging to Class XII. This became known as the Shapley–Sawyer Concentration Class. (It is sometimes given with numbers [Class 1–12] rather than Roman numerals.)^[16]

Formation

NGC 2808 contains three distinct generations of stars.^[17] NASA image

At present, the formation of globular clusters remains a poorly understood phenomenon, and it remains uncertain whether the stars in a globular cluster form in a single generation, or are spawned across multiple generations over a period of several hundred million years. In many globular clusters, most of the stars are at approximately the same stage in stellar evolution, suggesting that they formed at about the same time.^[18] However, the star formation history varies from cluster to cluster, with some clusters showing distinct populations of stars. An example of this is the globular clusters in the Large Magellanic Cloud (LMC) that exhibit a bimodal population. During their youth, these LMC clusters may have encountered giant molecular clouds that triggered a second round of star formation.^[19] This star-forming period is relatively brief, compared to the age of many globular clusters.^[20] It has also been proposed that the reason for this multiplicity in stellar populations could have a dynamical origin. In e.g. the Antennae galaxy we observe thanks to the Hubble Space Telescope clusters of clusters, regions in the galaxy that span hundreds of parsec, where many of the clusters will eventually collide and merge. Many of them present a significant range in ages, hence possibly metallicities, and their merger could plausibly lead to clusters with a bimodal or even multiple distribution of populations.^[21]

Globular star cluster Messier 54.^[22]

Observations of globular clusters show that these stellar formations arise primarily in regions of efficient star formation, and where the interstellar medium is at a higher density than in normal star-forming regions. Globular cluster formation is prevalent in starburst regions and in interacting galaxies.^[23] Research indicates a correlation between the mass of a central supermassive black holes (SMBH) and the extent of the globular cluster systems of elliptical and lenticular galaxies. The mass of the SMBH in such a galaxy is often close to the combined mass of the galaxy's globular clusters.^[24]

No known globular clusters display active star formation, which is consistent with the view that globular clusters are typically the oldest objects in the 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.^[25]

Composition

Djorgovski 1's stars contain hydrogen and helium, but not much else. In astronomical terms, they are described as "metal-poor".^[26]

Globular clusters are generally composed of hundreds of thousands of low-metal, old stars. The type of stars found in a globular cluster are similar to those in the bulge of a spiral galaxy but confined to a volume of only a few million cubic parsecs. They are free of gas and dust and it is presumed that all of the gas and dust was long ago turned into 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 per cubic parsec in the core of the cluster.^[27] The typical distance between stars in a globular cluster is about 1 light year,^[28] but at its core, the separation is comparable to the size of the Solar System (100 to 1000 times closer than stars near the Solar System).^[29]

However, they are not thought to be favorable locations for the survival of planetary systems. Planetary orbits are dynamically unstable within the cores of dense clusters because of the perturbations of passing stars. A planet orbiting at 1 astronomical unit around a star that is within the core of a dense cluster such as 47 Tucanae would only survive on the order of 10⁸ years.^[30] 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.^[31]

Some globular clusters, like Omega Centauri in our Milky Way and G1 in M31, are extraordinarily massive, with several million solar masses and multiple stellar populations. Both can be regarded as evidence that supermassive globular clusters are in fact the cores of dwarf galaxies that are consumed by the larger galaxies.^[32] About a quarter of the globular cluster population in the Milky Way may have been accreted along with their host dwarf galaxy.^[33]

Several globular clusters (like M15) have extremely massive cores which may harbor black holes,^[34] although simulations suggest that a less massive black hole or central concentration of neutron stars or massive white dwarfs explain observations equally well.

Metallic content

Messier 53 has surprised astronomers with its unusual number of a type of star called blue stragglers.^[35]

Globular clusters normally consist of Population II stars, which have a low proportion of elements other than hydrogen and helium when compared to Population I stars such as the Sun. Astronomers refer to these heavier elements as metals and to the proportions of these elements as the metallicity. These elements are produced by stellar nucleosynthesis and then are recycled into the interstellar medium, where they enter the next generation of stars. Hence the proportion of metals can be an indication of the age of a star, with older stars typically having a lower metallicity.^[36]

The Dutch astronomer Pieter Oosterhoff noticed that there appear to be two populations of globular clusters, which became known as Oosterhoff groups. The second group has a slightly longer period of RR Lyrae variable stars.^[37] Both groups have weak lines of metallic elements. But the lines in the stars of Oosterhoff type I (OoI) cluster are not quite as weak as those in type II (OoII).^[37] Hence type I are referred to as "metal-rich" (e.g. Terzan 7^[38]) while type II are "metal-poor" (e.g. ESO 280-SC06^[39]).

These two 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, but differ in their metal abundances. Many scenarios have been suggested to explain these subpopulations, including violent gas-rich galaxy mergers, the accretion of dwarf galaxies, and multiple phases of star formation in a single galaxy. In our Milky Way, the metal-poor clusters are associated with the halo and the metal-rich clusters with the bulge.^[40]

In the Milky Way it has been discovered that the large majority of the low metallicity clusters are aligned along a plane in the outer part of the galaxy's halo. This result argues in favor of the view that type II clusters in the galaxy were captured from a satellite galaxy, rather than being the oldest members of the Milky Way's globular cluster system as had been 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.^[41]

Exotic components

Globular clusters have a very high star density, and therefore close interactions and near-collisions of stars occur relatively often. Due to these chance encounters, some exotic classes of stars, such as blue stragglers, millisecond pulsars and low-mass X-ray binaries, are much more common in globular clusters. A blue straggler is formed from the merger of two stars, possibly as a result of an encounter with a binary system.^[42] The resulting star has a higher temperature than comparable stars in the cluster with the same luminosity, and thus differs from the main sequence stars formed at the beginning of the cluster.^[43]

Globular cluster M15 may have an intermediate-mass black hole at its core. NASA image.

Astronomers have searched for black holes within globular clusters since the 1970s. The resolution requirements for this task, however, are exacting, and it is only with the Hubble space telescope that the first confirmed discoveries have been made. In independent programs, a 4,000 solar mass intermediate-mass black hole has been suggested to exist based on HST observations in the globular cluster M15 and a 20,000 solar mass black hole in the Mayall II cluster in the Andromeda Galaxy.^[44] Both x-ray and radio emissions from Mayall II appear to be consistent with an intermediate-mass black hole.^[45]

These are of particular interest because they are the first black holes discovered that were intermediate in mass between the conventional stellar-mass black hole and the supermassive black holes discovered at the cores of galaxies. The mass of these intermediate mass black holes is proportional to the mass of the clusters, following a pattern previously discovered between supermassive black holes and their surrounding galaxies.

Claims of intermediate mass black holes have been met with some skepticism. The heaviest objects in globular clusters are expected to migrate to the cluster center due to mass segregation. As pointed out in two papers by Holger Baumgardt and collaborators, the mass-to-light ratio should rise sharply towards the center of the cluster, even without a black hole, in both M15^[46] and Mayall II.^[47]

Color-magnitude diagram

Messier 5 is a globular cluster consisting of hundreds of thousands of stars bound together by their collective gravity.^[48]

The Hertzsprung-Russell diagram (HR-diagram) is a graph of a large sample of stars that plots their visual absolute magnitude against their color index. The color index, B−V, is the difference between the magnitude of the star in blue light, or B, and the magnitude in visual light (green-yellow), or V. Large positive values indicate a red star with a cool surface temperature, while negative values imply a blue star with a hotter surface.

When the stars near the Sun are plotted on an HR diagram, it displays a distribution of stars of various masses, ages, and compositions. Many of the stars lie relatively close to a sloping curve with increasing absolute magnitude as the stars are hotter, known as main-sequence stars. However the diagram also typically includes stars that are in later stages of their evolution and have wandered away from this main-sequence curve.

As all the stars of a globular cluster are at approximately the same distance from us, their absolute magnitudes differ from their visual magnitude by about the same amount. The main-sequence stars in the globular cluster will fall along a line that is believed to be comparable to similar stars in the solar neighborhood. The accuracy of this assumption is confirmed by comparable results obtained by comparing the magnitudes of nearby short-period variables, such as RR Lyrae stars and cepheid variables, with those in the cluster.^[49]

By matching up these curves on the HR diagram the absolute magnitude of main-sequence stars in the cluster can also be determined. This in turn provides a distance estimate to the cluster, based on the visual magnitude of the stars. The difference between the relative and absolute magnitude, the distance modulus, yields this estimate of the distance.^[50]

When the stars of a particular globular cluster are plotted on an HR diagram, in many cases nearly all of the stars fall upon a relatively well defined curve. This differs from the HR diagram of stars near the Sun, which lumps together stars of differing ages and origins. The shape of the curve for a globular cluster is characteristic of a grouping of stars that were formed at approximately the same time and from the same materials, differing only in their initial mass. As the position of each star in the HR diagram varies with age, the shape of the curve for a globular cluster can be used to measure the overall age of the star population.^[51]

Color-magnitude diagram for the globular cluster M3. Note the characteristic "knee" in the curve at magnitude 19 where stars begin entering the giant stage of their evolutionary path.

The most massive main-sequence stars will also have the highest absolute magnitude, and these will be the first to evolve into the giant star stage. As the cluster ages, stars of successively lower masses will also enter the giant star stage. Thus the age of a single population cluster can be measured by looking for the stars that are just beginning to enter the giant star stage. This forms a "knee" in the HR diagram, bending to the upper right from the main-sequence line. The absolute magnitude at this bend is directly a function of the age of globular cluster, so an age scale can be plotted on an axis parallel to the magnitude.

In addition, globular clusters can be dated by looking at the temperatures of the coolest white dwarfs. Typical results for globular clusters are that they may be as old as 12.7 billion years.^[52] This is in contrast to open clusters which are only tens of millions of years old.

The ages of globular clusters place a bound on the age limit of the entire universe. This lower limit has been a significant constraint in cosmology. During the early 1990s, astronomers were faced with age estimates of globular clusters that appeared older than cosmological models would allow. However, better measurements of cosmological parameters through deep sky surveys and satellites such as COBE have resolved this issue as have computer models of stellar evolution that have different models of mixing.^{[citation needed]}

Evolutionary studies of globular clusters can also be used to determine changes due to the starting composition of the gas and dust that formed the cluster. That is, the evolutionary tracks change with changes in the abundance of heavy elements. The data obtained from studies of globular clusters are then used to study the evolution of the Milky Way as a whole.^[53]

In globular clusters a few stars known as blue stragglers are observed, apparently continuing the main sequence in the direction of brighter, bluer stars. The origins of these stars is still unclear, but most models suggest that these stars are the result of mass transfer in multiple star systems.^[54]

Morphology

NGC 411 is classified as an open cluster.^[55]

In contrast to open clusters, most globular clusters remain gravitationally bound for time periods comparable to the life spans of the majority of their stars. However, a possible exception is when strong tidal interactions with other large masses result in the dispersal of the stars.

After they are formed, the stars in the globular cluster begin to interact gravitationally with each other. As a result the velocity vectors of the stars are steadily modified, and the stars lose any history of their original velocity. The characteristic interval for this to occur is the relaxation time. This is related to the characteristic length of time a star needs to cross the cluster as well as the number of stellar masses in the system.^[56] The value of the relaxation time varies by cluster, but the mean value is on the order of 10⁹ years.

Although globular clusters generally appear spherical in form, ellipticities can occur due to 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.^[58]

Radii

Astronomers characterize the morphology of a globular cluster by means of standard radii. These are the core radius (r_c), the half-light radius (r_h) and the tidal radius (r_t). The overall luminosity of the cluster steadily decreases with distance from the core, and the core radius is the distance at which the apparent surface luminosity has dropped by half.^[59] A comparable quantity is the half-light radius, or the distance from the core within which half the total luminosity from the cluster is received. This is typically larger than the core radius.

Note that the half-light radius includes stars in the outer part of the cluster that happen to lie along the line of sight, so theorists will also use the half-mass radius (r_m)—the radius from the core that contains half the total mass of the cluster. When the half-mass radius of a cluster is small relative to the overall size, it has a dense core. An example of this is Messier 3 (M3), which has an overall visible dimension of about 18 arc minutes, but a half-mass radius of only 1.12 arc minutes.^[60]

Almost all globular clusters have a half-light radius of less than 10 pc, although there are well-established globular clusters with very large radii (i.e. NGC 2419 (R_h = 18 pc) and Palomar 14 (R_h = 25 pc)).^[11]

Finally the tidal radius 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 is about 38 arc minutes.

Mass segregation, luminosity and core collapse

In measuring the luminosity curve of a given globular cluster as a function of distance from the core, most clusters in the Milky Way increase steadily in luminosity as this distance decreases, up to a certain distance from the core, then the luminosity levels off. Typically this distance is about 1–2 parsecs from the core. However about 20% of the globular clusters have undergone a process termed "core collapse". In this type of cluster, the luminosity continues to increase steadily all the way to the core region.^[61] An example of a core-collapsed globular is M15.

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

Core-collapse is thought to occur when the more massive stars in a globular cluster encounter their less massive companions. Over time, dynamic processes cause individual stars to migrate from the center of the cluster to the outside. This results in a net loss of kinetic energy from the core region, leading the remaining stars grouped in the core region 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.^[62] (Note that a core collapse is not the only mechanism that can cause such a luminosity distribution; a massive black hole at the core can also result in a luminosity cusp.)^[63] Over a lengthy period of 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 the primordial supply of binaries is exhausted due to interactions can a deeper core collapse proceed.^[64][65] 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.^[66]

The different stages of core-collapse may be divided into three phases. During a globular cluster's adolescence, the process of core-collapse begins with stars near the core. However, the interactions between binary star systems prevents further collapse as the cluster approaches middle age. Finally, 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, which causes the cluster to re-expand. As the mean 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.^[67]

The Hubble Space Telescope has been used to provide 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 globular star cluster 47 Tucanae, which is made up of about 1 million stars, is one of the densest globular clusters in the Southern Hemisphere. This cluster was subjected to an intensive photographic survey, which allowed astronomers to track the motion of its stars. Precise velocities were obtained for nearly 15,000 stars in this cluster.^[68]

A 2008 study by John Fregeau of 13 globular clusters in the Milky Way shows that three of them have an unusually large number of X-ray sources, or X-ray binaries, suggesting the clusters are middle-aged. Previously, these globular clusters had been classified as being in old age because they had very tight concentrations of stars in their centers, another test of age used by astronomers. The implication is that most globular clusters, including the other ten studied by Fregeau, are not in middle age as previously thought, but are actually in 'adolescence'.^[69]

The overall luminosities of the globular clusters within the Milky Way and the Andromeda Galaxy can be modeled by means of a gaussian curve. This gaussian can be represented by means of 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 = −7.20 ± 0.13, σ = 1.1 ± 0.1 magnitudes.)^[70] The GCLF has also been used as a "standard candle" for measuring the distance to other galaxies, under the assumption that the globular clusters in remote galaxies follow the same principles as they do in the Milky Way.

N-body simulations

Computing the interactions between the stars within a globular cluster requires solving what is termed the N-body problem. That is, each of the stars within the cluster continually interacts with the other N−1 stars, where N is the total number of stars in the cluster. The naive CPU computational "cost" for a dynamic simulation increases in proportion to N³,^[71][72] so the potential computing requirements to accurately simulate such a cluster can be enormous.^[73] An efficient method of mathematically simulating the N-body dynamics of a globular cluster is done by subdividing into small volumes and velocity ranges, and using probabilities to describe the locations of the stars. The motions are then described by means of a formula called the Fokker-Planck equation. This can be solved by a simplified form of the equation, or by running Monte Carlo simulations and using random values. However 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.^[74]

The results of N-body simulations have shown that the stars can follow unusual paths through the cluster, often forming loops and often falling more directly toward the core than would a single star orbiting a central mass. In addition, due to interactions with other stars that result in an increase in velocity, some of the stars gain sufficient energy to escape the cluster. Over long periods of time this will result in a dissipation of the cluster, a process termed evaporation.^[75] The typical time scale for the evaporation of a globular cluster is 10¹⁰ years.^[56] In 2010 it became possible to directly compute, star by star, N-body simulations of a globular cluster over the course of its lifetime.^[76]

Binary stars form a significant portion of the total population of stellar systems, with up to half of all stars occurring in binary systems. 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.^[43]

The ultimate fate of a globular cluster must be either to accrete stars at its core, causing its steady contraction,^[77] or gradual shedding of stars from its outer layers.^[78]

Intermediate forms

Messier 10 is a ball of stars that lies about 15000 light-years from Earth, in the constellation of Ophiuchus.^[79]

The distinction between cluster types is not always clear-cut, and objects have been found that blur the lines between the categories. For example, BH 176 in the southern part of the Milky Way has properties of both an open and a globular cluster.^[80]

In 2005, astronomers discovered a completely new type of star cluster in the Andromeda Galaxy, which is, in several ways, very similar to globular clusters. The new-found clusters contain hundreds of thousands of stars, a similar number to that found in globular clusters. The clusters share other characteristics with globular clusters such as stellar populations and metallicity. What distinguishes them from the globular clusters is that they are much larger – several hundred light-years across – and hundreds of times less dense. The distances between the stars are, therefore, much greater within the newly discovered extended clusters. Parametrically, these clusters lie somewhere between a globular cluster and a dwarf spheroidal galaxy.^[81]

How these clusters are formed is not yet known, but their formation might well be related to that of globular clusters. Why M31 has such clusters, while the Milky Way does not, is not yet known. It is also unknown if any other galaxy contains these types of clusters, but it would be very unlikely that M31 is the sole galaxy with extended clusters.^[81]

Tidal encounters

When a globular cluster has a close encounter with a large mass, such as the core region of a galaxy, it undergoes a tidal interaction. The difference in the pull of gravity between the part of the cluster nearest the mass and the pull on the furthest part of the cluster results in a tidal force. A "tidal shock" occurs whenever the orbit of a cluster takes it through the plane of a galaxy.

As a result of a tidal shock, streams of stars can be pulled away from the cluster halo, leaving only the core part of the cluster. These tidal interaction effects create tails of stars that can extend up to several degrees of arc away from the cluster.^[82] These tails typically both precede and follow the cluster along its orbit. The tails can accumulate significant portions of the original mass of the cluster, and can form clumplike features.^[83]

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 out to distances of 13,000 light-years.^[84] Tidal interactions have stripped away much of the mass from Palomar 5, and further interactions as it passes through the galactic core are expected to transform it into a long stream of stars orbiting the Milky Way halo.

Tidal interactions add kinetic energy into a globular cluster, dramatically increasing the evaporation rate and shrinking the size of the cluster.^[56] Not only does tidal shock strip off the outer stars from a globular cluster, but the increased evaporation accelerates the process of core collapse. The same physical mechanism may be at work in Dwarf spheroidal galaxies such as the Sagittarius Dwarf, which appears to be undergoing tidal disruption due to its proximity to the Milky Way.

Orbits

There are many globular clusters with a retrograde orbit round the Milky Way Galaxy.^[85] A hypervelocity globular cluster was discovered around Messier 87 in 2014, having a velocity in excess of the escape velocity of M87.^[86]

Planets

In 2000, the results of a search for giant planets in the globular cluster 47 Tucanae were announced.
The lack of any successful discoveries suggests that the abundance of elements (other than hydrogen or helium) necessary to build these planets may need to be at least 40% of the abundance in the Sun. Terrestrial planets are built from heavier elements such as silicon, iron and magnesium. The very low abundance of these elements in globular clusters means that the member stars have a far lower likelihood of hosting Earth-mass planets, when compared to stars in the neighborhood of the Sun. Hence the halo region of the Milky Way galaxy, including globular cluster members, are unlikely to host habitable terrestrial planets.^[87]

In spite of the lower likelihood of giant planet formation, just such an object has been found in the globular cluster Messier 4. This planet was detected orbiting a pulsar in the binary star system PSR B1620-26. The eccentric and highly inclined orbit of the planet suggests it may have been formed around another star in the cluster, then was later "exchanged" into its current arrangement.^[88] The likelihood of close encounters between stars in a globular cluster can disrupt planetary systems, some of which break loose to become free floating planets. Even close orbiting planets can become disrupted, potentially leading to orbital decay and an increase in orbital eccentricity and tidal effects.^[89]