DEPT. OF PHYSICS, ST. MARY’S COLLEGE MANARCAUD

ASTROPHYSICS VI SEMESTER NUCLEAR PARTICLE AND ASTROPHYSICS

BY DR. RESMI G NAIR

[Type the abstract of the document here. The abstract is typically a short summary of the contents of the document. Type the abstract of the document here. The abstract is typically a short summary of the contents of the document.] CLASSIFICATION OF

Stars are classified by their spectra (the elements that they absorb) and their temperature. There are seven main types of stars. In order of decreasing temperature, O, B, A, F, G, K, and M. O and B stars are uncommon but very bright; M stars are common but dim. Electromagnetic radiation from the is analysed by splitting it with a prism or diffraction grating into a spectrum exhibiting the rainbow of colours interspersed with spectral lines. Each line indicates a particular chemical element or molecule, with the line strength indicating the abundance of that element. The strengths of the different spectral lines vary mainly due to the temperature of the photosphere, although in some cases there are true abundance differences. The spectral class of a star is a short code primarily summarizing the ionization state, giving an objective measure of the photosphere's temperature. Most stars are currently classified under the Morgan-Keenan (MK) system using the letters O, B, A, F, G, K, and M, a sequence from the hottest (O type) to the coolest (M type). Each letter class is then subdivided using a numeric digit with 0 being hottest and 9 being coolest (e.g. A8, A9, F0, and F1 form a sequence from hotter to cooler). The sequence has been expanded with classes for other stars and star-like objects that do not fit in the classical system, such as class D for white dwarfs and classes S and C for stars. In the MK system, a luminosity class is added to the spectral class using Roman numerals. This is based on the width of certain absorption lines in the star's spectrum, which vary with the density of the atmosphere and so distinguish giant stars from dwarfs. Luminosity class 0 or Ia+ is used for , class I for supergiants, class II for bright giants, class III for regular giants, class IV for sub-giants, class V for main- sequence stars, class sd (or VI) for sub-dwarfs, and class D (or VII) for white dwarfs. The full spectral class for the is then G2V, indicating a main-sequence star with a surface temperature around 5,800 K.

Main - Main- Main- seque seque Fraction Chroma sequence Effective Vega- nce nce Hydro of all Cla ticity [ luminosi temperatu relative chromati [5][6 radius [1][7] gen main- ss [1][2] [3][4][a] (D65) 1][7] [1][7] ty

re city lines sequence s ][3][b] (bolomet (solar (solar tars[8] ric) mass radii) es)

≥ 16 ≥ 6.6 ≥ 30,000 ~0.00003

O ≥ 30,000 K blue blue Weak M☉ R☉ L☉ %

deep 2.1– 25– 10,000– 1.8– Mediu

B blue white blue 16 M 30,000 L 0.13% 30,000 K 6.6 R☉ m white ☉ ☉

1.4– 7,500– blue 1.4–

A white 2.1 M 5–25 L☉ Strong 0.6% 10,000 K white 1.8 R☉ ☉ 1.04– 6,000– 1.15– Mediu

F yellow white white 1.4 M 1.5–5 L☉ 3% 7,500 K 1.4 R☉ m ☉

0.8– 0.96– 5,200– yellowis 0.6–1.5 L

G yellow 1.04 1.15 R Weak 7.6% 6,000 K h white ☉ M☉ ☉

pale 0.45– 0.7– 3,700– 0.08– Very

K light orange yellow 0.8 M 0.96 R 12.1% 5,200 K 0.6 L☉ weak orange ☉ ☉

light 0.08– 2,400– ≤ 0.7 Very

M orange red orange 0.45 ≤ 0.08 L☉ 76.45% 3,700 K R☉ weak red M☉

The spectral classes O through M, are subdivided by Arabic numerals (0–9), where 0 denotes the hottest stars of a given class. For example, A0 denotes the hottest stars in class A and A9 denotes the coolest ones. Fractional numbers are allowed; for example, the star Mu Normae is classified as O9.7. The Sun is classified as G2. Conventional colour descriptions are traditional in astronomy, and represent colours relative to the mean colour of an A class star, which is considered to be white. The apparent colour descriptions are what the observer would see if trying to describe the stars under a dark sky without aid to the eye, or with binoculars. However, most stars in the sky, except the brightest ones, appear white or bluish white to the unaided eye because they are too dim for colour vision to work. Red supergiants are cooler and redder than dwarfs of the same spectral type, and stars with particular spectral features such as carbon stars may be far redder than any black body. The fact that the Harvard classification of a star indicated its surface or photosphere temperature (or more precisely, its effective temperature) was not fully understood until after its development, though by the time the first Hertzsprung–Russell diagram was formulated (by 1914), this was generally suspected to be true. In the 1920s, the Indian physicist Meghnad Saha derived a theory of ionization by extending well-known ideas in physical chemistry pertaining to the dissociation of molecules to the ionization of atoms. First he applied it to the solar chromosphere, then to stellar spectra. Class O O-type stars are very hot and extremely luminous, with most of their radiated output in the ultraviolet range. These are the rarest of all main-sequence stars. About 1 in 3,000,000 (0.00003%) of the main-sequence stars in the solar neighbourhood are O-type stars. Some of the most massive stars lie within this spectral class. O-type stars frequently have complicated surroundings that make measurement of their spectra difficult. O-type spectra formerly were defined by the ratio of the strength of the He II λ 4541 relative to that of He I λ 4471, where λ is the wavelength, measured in angstroms.

Class B B-type stars are very luminous and blue. Their spectra have neutral lines, which are most prominent at the B2 subclass, and moderate lines. As O- and B-type stars are so energetic, they only live for a relatively short time. Thus, due to the low probability of kinematic interaction during their lifetime, they are unable to stray far from the area in which they formed, apart from runaway stars. The transition from class O to class B was originally defined to be the point at which the He II λ4541 disappears. Class A A-type stars are among the more common naked eye stars, and are white or bluish-white. They have strong hydrogen lines, at a maximum by A0, and also lines of ionized metals (Fe II, Mg II, Si II) at a maximum at A5. The presence of Ca II lines is notably strengthening by this point. About 1 in 160 (0.625%) of the main-sequence stars in the solar neighbourhood are A-type stars. Class F F-type stars have strengthening spectral lines H and K of Ca II. Neutral metals (Fe I, Cr I) beginning to gain on ionized metal lines by late F. Their spectra are characterized by the weaker hydrogen lines and ionized metals. Their colour is white. About 1 in 33 (3.03%) of the main-sequence stars in the solar neighbourhood are F-type stars. Class G G-type stars, including the Sun,[10] have prominent spectral lines H and K of Ca II, which are most pronounced at G2. They have even weaker hydrogen lines than F, but along with the ionized metals, they have neutral metals. There is a prominent spike in the G band of CH molecules. Class G main- sequence stars make up about 7.5%, nearly one in thirteen, of the main-sequence stars in the solar neighbourhood. Class G contains the "Yellow Evolutionary Void". Supergiant stars often swing between O or B (blue) and K or M (red). While they do this, they do not stay for long in the yellow supergiant G class, as this is an extremely unstable place for a supergiant to be. Class K K-type stars are orangish stars that are slightly cooler than the Sun. They make up about 12% of the main-sequence stars in the solar neighbourhood. There are also giant K-type stars, which range from hypergiants like RW Cephei, to giants and supergiants, such as Arcturus, whereas orange dwarfs, like Alpha Centauri B, are main-sequence stars. They have extremely weak hydrogen lines, if those are present at all, and mostly neutral metals (Mn I, Fe I, Si I). By late K, molecular bands of titanium oxide become present. There is a suggestion that K-spectrum stars may potentially increase the chances of life developing on orbiting planets that are within the habitable zone. Class M Class M stars are by far the most common. About 76% of the main-sequence stars in the solar neighbourhood are class M stars. However, class M main-sequence stars (red dwarfs) have such low luminosities that none are bright enough to be seen with the unaided eye, unless under exceptional conditions. The brightest known M-class main-sequence star is M0V Lacaille 8760, with magnitude 6.6 (the limiting magnitude for typical naked-eye visibility under good conditions is typically quoted as 6.5), and it is extremely unlikely that any brighter examples will be found. Although most class M stars are red dwarfs, most of the largest ever supergiant stars in the Milky Way are M stars, such as VV Cephei, Antares, and Betelgeuse, which are also class M. Furthermore, the larger, hotter brown dwarfs are late class M, usually in the range of M6.5 to M9.5. The spectrum of a class M star contains lines from oxide molecules (in the visible spectrum, especially TiO) and all neutral metals, but absorption lines of hydrogen are usually absent. TiO bands can be strong in class M stars, usually dominating their visible spectrum by about M5. Vanadium(II) oxide bands become present by late M.

Stellar evolution

Representative lifetimes of stars as a function of their

Stellar evolution is the process by which a star changes over the course of time. Depending on the mass of the star, its lifetime can range from a few million years for the most massive to trillions of years for the least massive, which is considerably longer than the age of the universe. The table shows the lifetimes of stars as a function of their masses. All stars are born from collapsing clouds of gas and dust, often called nebulae or molecular clouds. Over the course of millions of years, these protostars settle down into a state of equilibrium, becoming what is known as a main- sequence star. powers a star for most of its life. Initially the energy is generated by the fusion of hydrogen atoms at the core of the main-sequence star. Later, as the preponderance of atoms at the core becomes helium, stars like the Sun begin to fuse hydrogen along a spherical shell surrounding the core. This process causes the star to gradually grow in size, passing through the sub giant stage until it reaches the phase. Stars with at least half the mass of the Sun can also begin to generate energy through the fusion of helium at their core, whereas more-massive stars can fuse heavier elements along a series of concentric shells. Once a star like the Sun has exhausted its nuclear fuel, its core collapses into a dense and the outer layers are expelled as a . Stars with around ten or more times the mass of the Sun can explode in a as their inert iron cores collapse into an extremely dense or black hole. Although the universe is not old enough for any of the smallest red dwarfs to have reached the end of their lives, stellar models suggest they will slowly become brighter and hotter before running out of hydrogen fuel and becoming low-mass white dwarfs. Stellar evolution is not studied by observing the life of a single star, as most stellar changes occur too slowly to be detected, even over many centuries. Instead, astrophysicists come to understand how stars evolve by observing numerous stars at various points in their lifetime, and by simulating stellar structure using computer models. Protostar

Stellar evolution starts with the of a giant molecular cloud. Typical giant molecular clouds are roughly 100 light-years (9.5×1014 km) across and contain up to 6,000,000 solar masses (1.2×1037 kg). As it collapses, a giant molecular cloud breaks into smaller and smaller pieces. In each of these fragments, the collapsing gas releases gravitational potential energy as heat. As its temperature and pressure increase, a fragment condenses into a rotating sphere of superhot gas known as a protostar. A protostar continues to grow by accretion of gas and dust from the molecular cloud, becoming a pre- main-sequence star as it reaches its final mass. Further development is determined by its mass. Mass 30 is typically compared to the mass of the Sun: 1.0 M☉ (2.0×10 kg) means 1 . Protostars are encompassed in dust, and are thus more readily visible at infrared wavelengths. Brown dwarfs and sub-stellar objects 29 Protostars with masses less than roughly 0.08 M☉ (1.6×10 kg) never reach temperatures high enough for nuclear fusion of hydrogen to begin. These are known as brown dwarfs. The International Astronomical Union defines brown dwarfs as stars massive enough to fuse deuterium at some point in 28 their lives (13 Jupiter masses (MJ), 2.5 × 10 kg, or 0.0125 M☉). Objects smaller than 13 MJ are classified as sub-brown dwarfs (but if they orbit around another stellar object they are classified as planets). Both types, deuterium-burning and not, shine dimly and die away slowly, cooling gradually over hundreds of millions of years.

Main sequence For a more-massive protostar, the core temperature will eventually reach 10 million kelvin, initiating the proton–proton chain reaction and allowing hydrogen to fuse, first to deuterium and then to helium. 30 In stars of slightly over 1 M☉ (2.0×10 kg), the carbon–nitrogen–oxygen fusion reaction (CNO cycle) contributes a large portion of the energy generation. The onset of nuclear fusion leads relatively quickly to a hydrostatic equilibrium in which energy released by the core maintains a high gas pressure, balancing the weight of the star's matter and preventing further gravitational collapse. The star thus evolves rapidly to a stable state, beginning the main-sequence phase of its evolution. A new star will sit at a specific point on the of the Hertzsprung–Russell diagram, with the main-sequence spectral type depending upon the mass of the star. Small, relatively cold, low- mass red dwarfs fuse hydrogen slowly and will remain on the main sequence for hundreds of billions of years or longer, whereas massive, hot O-type stars will leave the main sequence after just a few million years. A mid-sized yellow dwarf star, like the Sun, will remain on the main sequence for about 10 billion years. The Sun is thought to be in the middle of its main sequence lifespan. Mature stars Eventually the core exhausts its supply of hydrogen and the star begins to evolve off of the main sequence. Without the outward pressure generated by the fusion of hydrogen to counteract the force of the core contracts until either electron degeneracy pressure becomes sufficient to oppose gravity or the core becomes hot enough (around 100 MK) for helium fusion to begin. Which of these happens first depends upon the star's mass.

Supernova When the core of a massive stars collapses, it will form a neutron star, or in the case of cores that exceed the Tolman-Oppenheimer-Volkoff limit, a black hole. Through a process that is not completely understood, some of the gravitational potential energy released by this core collapse is converted into a Type Ib, Type Ic, or Type II supernova. It is known that the core collapse produces a massive surge of neutrinos, as observed with supernova SN 1987A. The extremely energetic neutrinos fragment some nuclei; some of their energy is consumed in releasing nucleons, including neutrons, and some of their energy is transformed into heat and kinetic energy, thus augmenting the started by rebound of some of the infalling material from the collapse of the core. Electron capture in very dense parts of the infalling matter may produce additional neutrons. Because some of the rebounding matter is bombarded by the neutrons, some of its nuclei capture them, creating a spectrum of heavier-than-iron material including the radioactive elements up to (and likely beyond) uranium. Although non-exploding red giants can produce significant quantities of elements heavier than iron using neutrons released in side reactions of earlier nuclear reactions, the abundance of elements heavier than iron (and in particular, of certain isotopes of elements that have multiple stable or long-lived isotopes) produced in such reactions is quite different from that produced in a supernova. Neither abundance alone matches that found in the Solar System, so both supernovae and ejection of elements from red giants are required to explain the observed abundance of heavy elements and isotopes thereof. The energy transferred from collapse of the core to rebounding material not only generates heavy elements, but provides for their acceleration well beyond escape velocity, thus causing a Type Ib, Type Ic, or Type II supernova. Current understanding of this energy transfer is still not satisfactory; although current computer models of Type Ib, Type Ic, and Type II supernovae account for part of the energy transfer, they are not able to account for enough energy transfer to produce the observed ejection of material.[25] However, neutrino oscillations may play an important role in the energy transfer problem as they not only affect the energy available in a particular flavour of neutrinos but also through other general-relativistic effects on neutrinos. Some evidence gained from analysis of the mass and orbital parameters of binary neutron stars (which require two such supernovae) hints that the collapse of an oxygen-neon-magnesium core may produce a supernova that differs observably (in ways other than size) from a supernova produced by the collapse of an iron core. The most massive stars that exist today may be completely destroyed by a supernova with an energy greatly exceeding its gravitational binding energy. This rare event, caused by pair-instability, leaves behind no black hole remnant. In the past history of the universe, some stars were even larger than the largest that exists today, and they would immediately collapse into a black hole at the end of their lives, due to photodisintegration. Stellar remnants

White and

A white dwarf, also called a degenerate dwarf, is a stellar core remnant composed mostly of electron-. A white dwarf is very dense: its mass is comparable to that of the Sun, while its volume is comparable to that of Earth. A white dwarf's faint luminosity comes from the emission of stored thermal energy; no fusion takes place in a white dwarf. The nearest known white dwarf is Sirius B, at 8.6 light years, the smaller component of the Sirius . There are currently thought to be eight white dwarfs among the hundred star systems nearest the Sun.[2] The unusual faintness of white dwarfs was first recognized in 1910.

White dwarfs are thought to be the final evolutionary state of stars whose mass is not high enough to become a neutron star, that of about 10 solar masses. This includes over 97% of the other stars in the Milky Way. After the hydrogen-fusing period of a main-sequence star of low or medium mass ends, such a star will expand to a red giant during which it fuses helium to carbon and oxygen in its core by the triple-alpha process. If a red giant has insufficient mass to generate the core temperatures required to fuse carbon (around 1 billion K), an inert mass of carbon and oxygen will build up at its center. After such a star sheds its outer layers and forms a planetary nebula, it will leave behind a core, which is the remnant white dwarf. Usually, white dwarfs are composed of carbon and oxygen.

The material in a white dwarf no longer undergoes fusion reactions, so the star has no source of energy. As a result, it cannot support itself by the heat generated by fusion against gravitational collapse, but is supported only by electron degeneracy pressure, causing it to be extremely dense. The physics of degeneracy yields a maximum mass for a non-rotating white dwarf, the —approximately 1.44 times M☉—beyond which it cannot be supported by electron degeneracy pressure. A carbon-oxygen white dwarf that approaches this mass limit, typically by mass transfer from a companion star, may explode as a via a process known as carbon detonation.

A white dwarf is very hot when it forms, but because it has no source of energy, it will gradually cool as it radiates its energy. This means that its radiation, which initially has a high color temperature, will lessen and redden with time. Over a very long time, a white dwarf will cool and its material will begin to crystallize, starting with the core. The star's low temperature means it will no longer emit significant heat or light, and it will become a cold black dwarf.[5] Because the length of time it takes for a white dwarf to reach this state is calculated to be longer than the current age of the universe (approximately 13.8 billion years), it is thought that no black dwarfs yet exist. The oldest white dwarfs still radiate at temperatures of a few thousand kelvins.

Neutron stars A neutron star is the collapsed core of a giant star which before collapse had a total mass of between 10 and 29 solar masses. Neutron stars are the smallest and densest stars, excluding black holes, hypothetical white holes, quark stars and strange stars.[1] Neutron stars have a radius on the order of 10 kilometres (6.2 mi) and a mass of about 1.4 solar masses. They result 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. Once formed, they no longer actively generate heat, and cool over time; however, they may still evolve further through collision or accretion. Most of the basic models for these objects imply that neutron stars are composed almost entirely of neutrons (subatomic particles with no net electrical charge and with slightly larger mass than protons); the electrons and protons present in normal matter combine to produce neutrons at the conditions in a neutron star. Neutron stars are partially supported against further collapse by neutron degeneracy pressure, a phenomenon described by the Pauli exclusion principle, just as white dwarfs are supported against collapse by electron degeneracy pressure. However neutron degeneracy pressure is not by itself sufficient to hold up an object beyond

0.7M☉ and repulsive nuclear forces play a larger role in supporting more massive neutron stars. If the remnant star has a mass exceeding the Tolman–Oppenheimer–Volkoff limit of around 2 solar masses, the combination of degeneracy pressure and nuclear forces is insufficient to support the neutron star and it continues collapsing to form a black hole. Neutron stars that can be observed are very hot and typically have a surface temperature of around 600000 K. They are so dense that 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 kilometre chunk of the Earth (a cube with edges of about 800 metres) from Earth's surface. Their magnetic fields are between 108 and 1015 (100 million to 1 quadrillion) times stronger than Earth's magnetic field. The gravitational field at the neutron star's surface is about 2×1011 (200 billion) times that of Earth's gravitational field. Black hole

A black hole is a region of space-time exhibiting gravitational acceleration so strong that nothing— no particles or even electromagnetic radiation such as light—can escape from it. The theory of general relativity predicts that a sufficiently compact mass can deform space-time to form a black hole. The boundary of the region from which no escape is possible is called the event horizon. Although the event horizon has an enormous effect on the fate and circumstances of an object crossing it, no locally detectable features appear to be observed. In many ways, a black hole acts like an ideal black body, as it reflects no light. Moreover, quantum field theory in curved space-time predicts that event horizons emit Hawking radiation, with the same spectrum as a black body of a temperature inversely proportional to its mass. This temperature is on the order of billionths of a kelvin for black holes of stellar mass, making it essentially impossible to observe.

Black holes of stellar mass are expected to form when very massive stars collapse at the end of their life cycle. After a black hole has formed, it can continue to grow by absorbing mass from its surroundings. By absorbing other stars and merging with other black holes, supermassive black holes of millions of solar masses (M☉) may form. There is consensus that supermassive black holes exist in the centers of most galaxies

The presence of a black hole can be inferred through its interaction with other matter and with electromagnetic radiation such as visible light. Matter that falls onto a black hole can form an external accretion disk heated by friction, forming some of the brightest objects in the universe. If there are other stars orbiting a black hole, their orbits can be used to determine the black hole's mass and location. Such observations can be used to exclude possible alternatives such as neutron stars. In this way, astronomers have identified numerous candidates in binary systems, and established that the radio source known as Sagittarius A*, at the core of the Milk y Way galaxy, contains a supermassive black hole of about 4.3 million solar masses.

The supermassive black hole at the core of supergiant elliptical galaxy Messier 87, with a mass ~7 billion times the Sun's, as depicted in the first image released by the Event Horizon Telescope (10 April 2019). Visible are the crescent-shaped emission ring and central shadow, which are gravitationally magnified views of the black hole's photon ring and the photon capture zone of its event horizon. The crescent shape arises from the black hole's rotation and relativistic beaming; the shadow is about 2.6 times the diameter of the event horizon. Chandrasekhar limit

The Chandrasekhar limit is the maximum mass of a stable white dwarf star. The currently accepted 30 value of the Chandrasekhar limit is about 1.4 M☉ (2.765×10 kg). White dwarfs resist gravitational collapse primarily through electron degeneracy pressure. The Chandrasekhar limit is the mass above which electron degeneracy pressure in the star's core is insufficient to balance the star's own gravitational self-attraction. Consequently, a white dwarf with a mass greater than the limit is subject to further gravitational collapse, evolving into a different type of stellar remnant, such as a neutron star or black hole. Those with masses under the limit remain stable as white dwarfs Electron degeneracy pressure is a quantum-mechanical effect arising from the Pauli Exclusion Principle. Since electrons are fermions, no two electrons can be in the same state, so not all electrons can be in the minimum-energy level. Rather, electrons must occupy a band of energy levels. Compression of the electron gas increases the number of electrons in a given volume and raises the maximum energy level in the occupied band. Therefore, the energy of the electrons increases on compression, so pressure must be exerted on the electron gas to compress it, producing electron degeneracy pressure. With sufficient compression, electrons are forced into nuclei in the process of electron capture, relieving the pressure.

Calculated values for the limit vary depending on the nuclear composition of the mass. Chandrasekhar gives the following expression, based on the equation of state for an ideal Fermi gas:

where:

 ħ is the reduced Planck constant  c is the  G is the gravitational constant

 μe is the average molecular weight per electron, which depends upon the chemical composition of the star.

 mH is the mass of the hydrogen atom. 0  ω 3 ≈ 2.018236 is a constant connected with the solution to the Lane–Emden equation.

Photon diffusion is a situation where photons travel through a material without being absorbed, but rather undergoing repeated scattering events which change the direction of their path. The path of any given photon is then effectively a random walk. A large ensemble of such photons can be said to exhibit diffusion in the material, and can be described with a diffusion equation. In astrophysics, photon diffusion occurs inside a stellar atmosphere. Hertzsprung–Russell diagram

The Hertzsprung–Russell diagram, abbreviated as H–R diagram, HR diagram or HRD, is a scatter plot of stars showing the relationship between the stars' absolute magnitudes or luminosities versus their stellar classifications or effective temperatures. More simply, it plots a star's luminosity (brightness) against its temperature (colour). The diagram was created circa 1910 by Ejnar Hertzsprung and Henry Norris Russell and represents a major step towards an understanding of stellar evolution.

An HR diagram showing many well-known stars in the Milky Way galaxy

There are several forms of the Hertzsprung–Russell diagram, and the nomenclature is not very well defined. All forms share the same general layout: stars of greater luminosity are toward the top of the diagram, and stars with higher surface temperature are toward the left side of the diagram.

The original diagram displayed the spectral type of stars on the horizontal axis and the absolute visual magnitude on the vertical axis. The spectral type is not a numerical quantity, but the sequence of spectral types is a monotonic series that reflects the stellar surface temperature. Modern observational versions of the chart replace spectral type by a colour index (in diagrams made in the middle of the 20th Century, most often the B-V colour) of the stars. This type of diagram is what is often called an observational Hertzsprung–Russell diagram, or specifically a colour–magnitude diagram (CMD), and it is often used by observers. In cases where the stars are known to be at identical distances such as within a star cluster, a colour–magnitude diagram is often used to describe the stars of the cluster with a plot in which the vertical axis is the apparent magnitude of the stars. For cluster members, by assumption there is a single additive constant difference between their apparent and absolute magnitudes, called the distance modulus, for all of that cluster of stars. Early studies of nearby open clusters (like the Hyades and Pleiades) by Hertzsprung and Rosenberg produced the first CMDs, a few years before Russell's influential synthesis of the diagram collecting data for all stars for which absolute magnitudes could be determined.

Another form of the diagram plots the effective surface temperature of the star on one axis and the luminosity of the star on the other, almost invariably in a log-log plot. Theoretical calculations of stellar structure and the evolution of stars produce plots that match those from observations. This type of diagram could be called temperature-luminosity diagram, but this term is hardly ever used; when the distinction is made, this form is called the theoretical Hertzsprung–Russell diagram instead. A peculiar characteristic of this form of the H–R diagram is that the temperatures are plotted from high temperature to low temperature, which aids in comparing this form of the H–R diagram with the observational form.

Although the two types of diagrams are similar, astronomers make a sharp distinction between the two. The reason for this distinction is that the exact transformation from one to the other is not trivial. To go between effective temperature and colour requires a colour–temperature relation, and constructing that is difficult; it is known to be a function of stellar composition and can be affected by other factors like stellar rotation.

Contemplation of the diagram led astronomers to speculate that it might demonstrate stellar evolution, the main suggestion being that stars collapsed from red giants to dwarf stars, then moving down along the line of the main sequence in the course of their lifetimes. Stars were thought therefore to radiate energy by converting gravitational energy into radiation through the Kelvin–Helmholtz mechanism. This mechanism resulted in an age for the Sun of only tens of millions of years, creating a conflict over the age of the Solar System between astronomers, and biologists and geologists who had evidence that the Earth was far older than that. This conflict was only resolved in the 1930s when nuclear fusion was identified as the source of stellar energy.

Most of the stars occupy the region in the diagram along the line called the main sequence. During the stage of their lives in which stars are found on the main sequence line, they are fusing hydrogen in their cores. The next concentration of stars is on the horizontal branch (helium fusion in the core and hydrogen burning in a shell surrounding the core). Another prominent feature is the Hertzsprung gap located in the region between A5 and G0 spectral type and between +1 and −3 absolute magnitudes (i.e. between the top of the main sequence and the giants in the horizontal branch). RR Lyrae variable stars can be found in the left of this gap on a section of the diagram called the instability strip. Cepheid variables also fall on the instability strip, at higher luminosities. RR Lyrae variables are periodic variable stars, commonly found in globular clusters. They are used as standard candles to measure (extra) galactic distances, assisting with the cosmic distance ladder. This class is named after the prototype and brightest example, RR Lyrae. A Cepheid variable is a type of star that pulsates radially, varying in both diameter and temperature and producing changes in brightness with a well-defined stable period and amplitude.

An HR diagram with the instability strip and its components highlighted