Variable Stars
Overview
Variable stars are stars that vary in luminosity. They do so because they are in an unstable and, for some, unpredictable part of their lifecycle. The specific cause of the variation is different for each type of variable star. This page explains the five main types of known variables.
T Tauri
T Tauri stars are named after the prototype t Tau, a relatively dim star in the constellation Taurus, the Bull. It is a star that is a member of the Hyades open cluster, and hence it is a very young star (about 1 million years old) and lies approximately 580 light-years from Earth. All stars like it are categorized as T Tauri stars.
T Tauris are young stars - like the prototype. They are usually still found with some of their original nebular material surrounding them. They are generally less than 2 solar masses and are pre-main sequence stars. These stars are still in the process of forming. It is thought that our own star had a T Tauri phase early in its history.
The cores of the stars are too low-temperature for fusion to take place, and so they produce energy by gravitational contraction. Theoretical estimates are that it takes approximately 100 million years for them to reach the main sequence and begin hydrogen fusion in their cores.
Their variability is non-regular and hence unpredictable, though it is generally over only about 0.5-1 magnitude. It is believed that the variability is due to clumps of material within their protoplanetary disk either eclipsing the star or falling into the star (around half of all T Tauri stars have these disks). Another source could be large starspots that travel over their surface, temporarily decreasing the luminosity.
Cepheid
Cepheids are named after the prototype delta (δ) Cep, the fourth-brightest star in the constellation Cepheus, the King. It is an old star for its mass, being about 100 million years in age. It lies approximately 891 light-years away. It weighs in at about 4-5 solar masses. All stars like it are categorized as Cepheid variable stars.
Cepheids are old stars for their giant masses - like the prototype. They are generally between 3-30 solar masses in weight and lie within the instability strip on the H-R diagram. These stars have reached a time in their lives when they have become unstable and will soon die. They are very bright stars, with, for example, δ Cep having a luminosity 500-2000 times the sun.
The instability is caused by a cycle of ionization of helium in the stars' atmosphere. As the ionization builds, the outer atmosphere becomes more opaque and the energy inside cannot get out. This puts pressure on it and the star expands and at the same time, the helium deionizes. The pulsation typically lasts a few days for a complete cycle, and the star's luminosity can vary by up to about 2 magnitudes.
The period of pulsation is directly linked to the star's luminosity, and hence they are very useful as indicators of distance, forming a foundational rung on the astronomical distance ladder. The period-luminosity relationship was first discovered by Henrietta Leavitt in 1912 (Leavitt & Pickering, 1912). She measured the brightness of hundreds of Cepheids and discovered that out of any two Cepheids with the same period of variation, the one with the brighter average magnitude is closer to us.
One of the recent calibrations comes from Hubble Space Telescope observations of Cepheids (Fritz et al., 2007), and it has calculated the relationship to be:
where MV is the absolute magnitude in the V-band Johnson filter and P is the period of the luminosity variation, measured in days.
RR Lyrae
RR Lyrae stars are named after the prototype RR Lyr, a faint star in the constellation Lyra. It is an old star and lies approximately 854 light-years away. All stars like it are categorized as RR Lyrae, though since this is a sub-type of Cepheids, they may be cross-listed.
The difference between RR Lyrae and Cepheids is that the period of time it takes for them to brighten and darken is measured on the order of many hours instead of many days. They are also fainter than normal Cepheids, with, for example, RR Lyr being approximately 50 times the luminosity of the sun. This is a consequence of their lower masses, typically about 0.8 solar masses.
Since they are a sub-type of Cepheids, they are also useful in the astronomical distance ladder, but because they are fainter, they cannot be used to determine the distances to other galaxies. One of the main uses is to be a calibration for globular cluster distances for use with the globular cluster luminosity function.
Mira
Mira variable stars are named after the prototype Mira, also known as omicron (ο) Cet, the 15th brightest star in the constellation Cetus, the Whale. Mira is a red giant star with an estimated age of 6 billion years. Its true distance is uncertain, with estimates ranging from 200-400 light-years. It weighs approximately 1.2 solar masses but has a luminosity of about 8500-9500 times that of the sun. It is also a member of a binary system. All stars that vary like it are known as Mira variables.
All Mira variables are old, red giant stars that may be indicative of what our sun may become, though their relatively higher mass may rule this out. Their pulsation periods range anywhere from 100 to 1000 days, and they can change in luminosity by up to a few magnitudes. They are unstable stars and will soon die, releasing a planetary nebula and becoming a white dwarf.
As a red giant star, they are no larger than about 2 solar masses, but their luminosity can be thousands of times that of the sun because of their extended outer atmospheres. Their change in brightness is believed to be a result of the entire star physically growing and shrinking in size, changing temperature and brightness while doing so. When the star grows in size, the temperature and luminosity drop because the same amount of energy is spread over a larger volume. When it contracts, the opposite occurs.
Unlike Cepheids and their cousins, RR Lyrae stars, Mira variables are a diverse group of stars that do not have any specific period-luminosity relationship. An individual star may also change its pulsation period from one cycle to the next, and it may change in brightness from one cycle to the next. Because of their slower periods and dramatic changes in brightness, they are a popular target for amateur astronomers interested in variable stars.
Cataclysmic Variable (Nova)
Cataclysmic variable stars are also known as nova (singular) or novae (plural). They should not be confused with supernovae events, where an entire star either collapses or explodes. These are events that require a binary star system with a white dwarf as the primary star and a still-normal secondary star from which mass is being transferred.
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| A not-to-scale illustration of a cataclysmic variable star system. Material from the giant star on the left is gravitationally accreted onto the surface of the smaller white dwarf on the right. Image created by Stuart Robbins. |
The stars in the binary system are so close that their orbital period is generally between 1-10 hours. The white dwarf can hence draw material from the secondary star companion, forming an accretion disk around it, and depositing material on its surface. Since stars are mostly hydrogen, most of the material drawn onto the surface of the white dwarf is hydrogen.
When enough material accumulates, the temperature and pressure on the bottom layer of hydrogen - in contact with the surface of the white dwarf - is enough to ignite in a brief burst of nuclear fusion, rapidly converting the hydrogen to helium. This brief burst of fusion is enough to change the brightness of the system by up to 10 magnitudes, making them obviously visible if in the nearby universe.
The brief increase in brightness is followed by a slow (several-day) decrease to its previous luminosity. These have characteristic timescales linked to their maximum luminosity, and hence they are useful in the astronomical distance ladder.
After the white dwarf returns to its quiescent state, it continues to draw material from the companion star. However, it can take centuries or millennia before it has accumulated enough to ignite again. If the star accumulates enough material that it weighs more than 1.4 solar masses (the Chandrasekhar limit), then the white dwarf will collapse into a neutron star.