The main sequence is the name for a continuous and distinctive band of stars that appear on a plot of stellar color versus brightness. These color-magnitude plots are known as Hertzsprung-Russell diagrams after their co-developers, Ejnar Hertzsprung and Henry Norris Russell. Stars on this band are known as main-sequence stars or dwarf stars.
After a star has formed, it generates energy at the hot, dense core region through the nuclear fusion of hydrogen atoms into helium. During this stage of the star's lifetime, it is located along the main sequence at a position determined primarily by its mass, but also based upon its chemical composition and other factors. In general, the more massive the star the shorter its lifespan on the main sequence. After the hydrogen fuel at the core has been consumed, the star evolves away from the main sequence.
The main sequence is sometimes divided into upper and lower parts, based on the processes that stars use to generate energy. Stars below about 1.5 times the mass of the Sun (or 1.5 solar masses) fuse hydrogen atoms together in a series of stages to form helium; a sequence called the proton-proton chain. Above this mass, in the upper main sequence, the nuclear fusion process can instead use atoms of carbon, nitrogen and oxygen as intermediaries in the production of helium from hydrogen atoms.
Because there is a temperature gradient between the core of a star and its surface, energy is steadily transported upward through the intervening layers until it is radiated away at the photosphere. The two mechanisms used to carry this energy through the star are radiation and convection, with the type used depending on the local conditions. Convection tends to occur in regions with steeper temperature gradients, higher opacity or both. When convection occurs in the core region it acts to stir up the helium ashes, thus maintaining the proportion of fuel needed for fusion to occur.
In the early part of the twentieth century, information about the types and distances of stars became more readily available. The spectra of stars were shown to have distinctive features, which allowed them to be categorized. Annie Jump Cannon and Edward C. Pickering at Harvard College Observatory had developed a method of categorization that became known as the Harvard classification scheme. This scheme was published in the Harvard Annals in 1901.
In Potsdam in 1906, the Danish astronomer Ejnar Hertzsprung noticed that the reddest stars—classified as K and M in the Harvard scheme—could be divided into two distinct groups. These stars are either much brighter than the Sun, or much fainter. To distinguish these groups, he called them "giant" and "dwarf" stars. The following year he began studying star clusters; large groupings of stars that are co-located at approximately the same distance. He published the first plots of color versus luminosity for these stars. These plots showed a prominent and continuous sequence of stars, which he named the main sequence.
At Princeton University, Henry Norris Russell was following a similar course of research. He was studying the relationship between the spectral classification of stars and their actual brightness as corrected for distance—their absolute magnitude. For this purpose he used a set of stars that had reliable parallaxes and many of which had been categorized at Harvard. When he plotted the spectral types of these stars against their absolute magnitude, he found that dwarf stars followed a distinct relationship. This allowed the real brightness of a dwarf star to be predicted with reasonable accuracy.
Of the red stars observed by Hertzsprung, the dwarf stars also followed the spectra-luminosity relationship discovered by Russell. However, the giant stars are much brighter than dwarfs and so do not follow the same relationship. Russell proposed that the "giant stars must have low density or great surface-brightness, and the reverse is true of dwarf stars". The same curve also showed that there were very few faint white stars.
In 1933, Bengt Strömgren introduced the term Hertzsprung-Russell diagram to denote a luminosity-spectral class diagram. This name reflected the parallel development of this technique by both Hertzsprung and Russell earlier in the century.
As evolutionary models of stars were developed during the 1930s, it was shown that, for stars of a uniform chemical composition, a relationship exists between a star's mass, composition, luminosity and radius. That is, the radius and luminosity of a star can be calculated given its mass and composition. This became known as the Vogt-Russell theorem; named after Heinrich Vogt and Henry Norris Russell. By this theorem, once a star's chemical composition and its position on the main sequence is known, so too is the star's mass and radius. (However, it was subsequently discovered that the theorem breaks down somewhat for stars of non-uniform composition.)
A refined scheme for stellar classification was published in 1943 by W. W. Morgan and P. C. Keenan. The MK classification assigned each star a spectral type—based on the Harvard classification—and a luminosity class. For historical reasons, the spectral types of stars followed, in order of decreasing temperature with colors ranging from blue to red, the sequence O, B, A, F, G, K and M. (A popular mnemonic for memorizing this sequence of stellar classes is "Oh Be A Fine Girl/Guy, Kiss Me".) The luminosity class ranged from I to V, in order of decreasing luminosity. Stars of luminosity class V belonged to the main sequence.
Main sequence stars have been extensively studied through stellar models, allowing their formation and evolutionary history to be relatively well understood. The position of the star on the main sequence provides information about its physical properties.
The temperature of a star can be approximately determined by treating it as an idealized energy radiator known as a black body. In this case, the luminosity L and radius R are related to the temperature T by the Stefan-Boltzmann Law:
where σ is the Stefan–Boltzmann constant. The temperature and composition of a star's photosphere determines the energy emission at different wavelengths. The color index, or B − V, measures the difference in this energy emission by means of filters that capture the star's magnitude in blue (B) and green-yellow (V) light. (By measuring the difference between these values, this eliminates the need to correct the magnitudes for distance.) Thus the position of a star on the HR diagram can be used to estimate its radius and temperature. By modifying the physical properties of the plasma in the photosphere, the temperature of a star also determines its spectral type.
When a protostar is formed from the collapse of a giant molecular cloud of gas and dust in the local interstellar medium, the initial composition is homogeneous throughout, consisting of about 70% hydrogen, 28% helium and trace amounts of other elements, by mass. During the initial collapse, this pre-main sequence star generates energy through gravitational contraction. Upon reaching a suitable density, energy generation is begun at the core using an exothermic nuclear fusion process that converts hydrogen into helium.
Once nuclear fusion of hydrogen becomes the dominant energy production process and the excess energy gained from gravitational contraction has been lost, the star lies along a curve on the Hertzsprung-Russell diagram (or HR diagram) called the standard main sequence. Astronomers will sometimes refer to this stage as "zero age main sequence", or ZAMS. This curve is calculated using computer models of stellar properties at the point when stars begin hydrogen fusion; the brightness and surface temperature of stars typically increase from this point with age.
A star remains near its initial position on the main sequence until a significant amount of hydrogen in the core has been consumed, then begins to evolve into a more luminous star. (On the HR diagram, the evolving star moves up and to the right of the main sequence.) Thus the main sequence represents the primary hydrogen-burning stage of a star's lifetime.
The majority of stars on a typical HR diagram lie along the main sequence curve. This line is so pronounced because both the spectral type and the luminosity depend only on a star's mass, at least to zeroth order approximation, as long as it is fusing hydrogen at its core—and that is what almost all stars spend most of their "active" life doing. These main-sequence (and therefore "normal") stars are called dwarf stars. This is not because they are unusually small, but instead comes from their smaller radii and lower luminosity as compared to the other main category of stars, the giant stars. White dwarfs are a different kind of star that are much smaller than main sequence stars—being roughly the size of the Earth. These represent the final evolutionary stage of many main sequence stars.
All main sequence stars have a core region where energy is generated by nuclear fusion. The temperature and density of this core are at the levels necessary to sustain the energy production that will support the remainder of the star. A reduction of energy production would cause the overlaying mass to compress the core, resulting in an increase in the fusion rate because of higher temperature and pressure. Likewise an increase in energy production would cause the star to expand, lowering the pressure at the core. Thus the star forms a self-regulating system in hydrostatic equilibrium that is stable over the course of its main sequence lifetime.
Astronomers divide the main sequence into upper and lower parts, based on the dominant type of fusion process at the core. Stars in the upper main sequence have sufficient mass to use the CNO cycle to fuse hydrogen into helium. This process uses atoms of carbon, nitrogen and oxygen as intermediaries in the fusion process. In the lower main sequence, energy is generated as the result of the proton-proton chain, which directly fuses hydrogen together in a series of stages to produce helium.
At a stellar core temperature of 18 million kelvins, both fusion processes are equally efficient. This is the core temperature of a star with 1.5 solar masses. Hence the upper main sequence consists of stars above this mass. The apparent upper limit for a main sequence star is 120-200 solar masses. Stars above this mass can not radiate energy fast enough to remain stable, so any additional mass will be ejected in a series of pulsations until the star reaches a stable limit. The lower limit for sustained nuclear fusion is about 0.08 solar masses.
Because there is a temperature difference between the core and the surface, or photosphere, energy is transported outward. The two modes for transporting this energy are radiation and convection. A radiation zone, where energy is transported by radiation, is stable against convection and there is very little mixing of the plasma. By contrast, in a convection zone the energy is transported by bulk movement of plasma, with hotter material rising and cooler material descending. Convection is a more efficient mode for carrying energy than radiation, but it will only occur under conditions that create a steep temperature gradient.
In massive stars, the rate of energy generation by the CNO cycle is very sensitive to temperature, so the fusion is highly concentrated at the core. Consequently, there is a high temperature gradient in the core region, which results in a convection zone for more efficient energy transport. This mixing of material around the core removes the helium ash from the hydrogen burning region, allowing more of the hydrogen in the star to be consumed during the main sequence lifetime. The outer regions of a massive star transport energy by radiation, with little or no convection.
Intermediate mass, class A stars such as Sirius may transport energy entirely by radiation. Medium-sized, low mass stars like the Sun have a core region that is stable against convection, with a convection zone near the surface. This produces mixing of the outer layers, but results in a less efficient consumption of the hydrogen in the star. This causes a steady buildup of a helium-rich core, surrounded by a hydrogen-rich outer region. By contrast, cool, low-mass stars are convective throughout. Thus the helium produced at the core is distributed across the star, producing a relatively uniform atmosphere and a proportionately longer main sequence lifespan.
As non-fusing helium ash accumulates in the core, the reduction in the abundance of hydrogen per unit mass results in a gradual lowering of the fusion rate within that mass. To compensate, the core temperature and pressure slowly increase, which actually causes a net increase in the overall fusion rate (to support the greater density of the inner star). This produces a steady increase in the luminosity and radius of the star over time. Thus, for example, the luminosity of the early Sun was only about 70% of its current value. As a star ages, the luminosity increase changes its position on the HR diagram. This effect results in a broadening of the main sequence band because stars are observed at random stages in their lifetime.
Other factors that broaden the main sequence band on the HR diagram include uncertainty in the distance to the stars, and the presence of unresolved binary stars that can alter the observed stellar parameters. However, even perfect observation would show a fuzzy main sequence, because mass is not the only parameter that affects a star's color and luminosity. In addition to variations in chemical composition—both because of the initial abundances and the star's evolutionary status, interaction with a close companion, rapid rotation, or a magnetic field can also change a main sequence star's position slightly on the HR diagram, to name just a few factors. For example, there are stars that have a very low abundance of elements with higher atomic numbers than helium—known as metal-poor stars—that lie just below the main sequence. Known as subdwarfs, these stars are also fusing hydrogen in their core and so they mark the lower edge of the main sequence's fuzziness due to chemical composition.
A nearly vertical region of the HR diagram, known as the instability strip, is occupied by pulsating variable stars. These stars vary in magnitude at regular intervals, giving them a pulsating appearance. The strip intersects the upper part of the main sequence in the region of class A and F stars; which are between one and two solar masses. However, main sequence stars in this region experience only small variations in magnitude and so are hard to detect.
The lifespan that a star spends on the main sequence is governed by two factors. The total amount of energy that can be generated through nuclear fusion of hydrogen is limited by the amount of available hydrogen fuel that can be consumed at the core. For a star in equilibrium, the energy generated at the core must be at least equal to the energy radiated at the surface. Since the luminosity gives the amount of energy radiated per unit time, the total life span can be estimated, to first approximation, as the total energy produced divided by the star's luminosity.
Our Sun has been a main sequence star for about 4.5 billion years and will continue to be one for another 5.5 billion years, for a total main sequence lifetime of 1010 years. After the hydrogen supply in the core is exhausted, it will expand to become a red giant and fuse helium atoms to form carbon. As the energy output of the helium fusion process per unit mass is only about a tenth the energy output of the hydrogen process, this stage will only last for about 10% of a star's total active lifetime. Thus, on average, about 90% of the observed stars will be on the main sequence.
The amount of fuel available for nuclear fusion is proportional to the mass of the star. Thus, the lifetime of a star on the main sequence can be estimated by comparing it to the Sun:
where M and L are the mass and luminosity of the star, respectively, is a solar mass, is the solar luminosity and is the star's estimated main sequence lifetime.
This is a counter-intuitive result, as more massive stars have more fuel to burn and might be expected to last longer. Instead, the lightest stars, of less than a tenth of a solar mass, may last over a trillion years. For the heaviest stars, however, this mass-luminosity relationship poorly matches the estimated lifetime, which last at least a few million years. A more accurate representation gives a different function for various ranges of mass.
The mass-luminosity relationship depends on how efficiently energy can be transported from the core to the surface. A higher opacity has an insulating effect that retains more energy at the core, so the star does not need to produce as much energy to remain in hydrostatic equilibrium. By contrast, a lower opacity means energy escapes more rapidly and the star must burn more fuel to remain in equilibrium. Note, however, that a sufficiently high opacity can result in energy transport via convection, which changes the conditions needed to remain in equilibrium.
In high mass main sequence stars, the opacity is dominated by electron scattering, which is nearly constant with increasing temperature. Thus the luminosity only increases as the cube of the star's mass. For stars below 10 times the solar mass, the opacity becomes dependent on temperature, resulting in the luminosity varying approximately as the fourth power of the star's mass. For very low mass stars, molecules in the atmosphere also contribute to the opacity. Below about 0.5 solar masses, the luminosity of the star varies as the mass to the power of 2.3, producing a flattening of the slope on a graph of mass versus luminosity. Even these refinements are only an approximation, however, and the mass-luminosity relation can vary depending on a star's composition.
Once a main sequence star consumes the hydrogen at its core, the loss of energy generation causes gravitational collapse to resume. The hydrogen surrounding the core reaches sufficient temperature and pressure to undergo fusion, forming a hydrogen-burning shell surrounding a helium core. In consequence of this change, the outer envelope of the star expands and decreases in temperature, turning it into a red giant. At this point the star is evolving off the main sequence and entering the giant branch. (The path the star now follows across the HR diagram is called an evolutionary track.) The helium core of the star continues to collapse until it is entirely supported by electron degeneracy pressure—a quantum mechanical effect that restricts how closely matter can be compacted. For stars of more than about 0.5 solar masses, the core can reach a temperature where it becomes hot enough to burn helium into carbon via the triple alpha process.
When a cluster of stars is formed at about the same time, the life span of these stars will depend on their individual masses. The most massive stars will leave the main sequence first, followed steadily in sequence by stars of ever lower masses. Thus the stars will evolve in order of their position on the main sequence, proceeding from the most massive at the left toward the right of the HR diagram. The current position where stars in this cluster are leaving the main sequence is known as the turn-off point. By knowing the main sequence lifespan of stars at this point, it becomes possible to estimate the age of the cluster.
The table below shows typical values for stars along the main sequence. The values of luminosity (L), radius (R) and mass (M) are relative to the Sun—a dwarf star with a spectral classification of G2 V. The actual values for a star may vary by as much as 20-30% from the values listed below. The coloration of the stellar class column gives an approximate representation of the star's photographic color, which is a function of the effective surface temperature.
Table of main sequence stellar parameters Stellar
Radius Mass Luminosity Temperature Examples R/R☉ M/M☉ L/L☉ K O5 18 40 500,000 38,000 Zeta Puppis B0 7.4 18 20,000 30,000 Phi1 Orionis B5 3.8 6.5 800 16,400 Pi Andromedae A A0 2.5 3.2 80 10,800 Alpha Coronae Borealis A A5 1.7 2.1 20 8,620 Beta Pictoris F0 1.4 1.7 6 7,240 Gamma Virginis F5 1.2 1.29 2.5 6,540 Eta Arietis G0 1.05 1.10 1.26 6,000 Beta Comae Berenices G2 1.00 1.00 1.00 5,920 Sun G5 0.93 0.93 0.79 5,610 Alpha Mensae K0 0.85 0.78 0.40 5,150 70 Ophiuchi A K5 0.74 0.69 0.16 — 61 Cygni A M0 0.63 0.47 0.063 3,920 Gliese 185 M5 0.32 0.21 0.0079 3,120 EZ Aquarii A M8 0.13 0.10 0.0008 — Van Biesbroeck's star
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