The radial velocity of an object with respect to a given point is the rate of change of the distance between the object and the point. That is, the radial velocity is the component of the object's velocity that points in the direction of the radius connecting the object and the point. In astronomy, the point is taken to be the observer on Earth, so the radial velocity denotes the speed with which the object moves away from or approaches the Earth. In astronomy, radial velocity is measured to the first order of approximation by Doppler spectroscopy; the quantity obtained by this method may be called the barycentric radial-velocity measure or spectroscopic radial velocity. However, due to relativistic and cosmological effects over the great distances that light travels to reach the observer from an astronomical object, this measure cannot be transformed to a geometric radial velocity without additional assumptions about the object and the space between it and the observer. By contrast, astrometric radial velocity is determined by astrometric observations.
Light from an object with a substantial relative radial velocity at emission will be subject to the Doppler effect, so the frequency of the light decreases for objects that were receding and increases for objects that were approaching. The radial velocity of a star or other luminous distant objects can be measured by taking a high-resolution spectrum and comparing the measured wavelengths of known spectral lines to wavelengths from laboratory measurements. A positive radial velocity indicates the distance between the objects was increasing. In many binary stars, the orbital motion causes radial velocity variations of several kilometers per second; as the spectra of these stars vary due to the Doppler effect, they are called spectroscopic binaries. Radial velocity can be used to estimate the ratio of the masses of the stars, some orbital elements, such as eccentricity and semimajor axis; the same method has been used to detect planets around stars, in the way that the movement's measurement determines the planet's orbital period, while the resulting radial-velocity amplitude allows the calculation of the lower bound on a planet's mass using the binary mass function.
Radial velocity methods alone may only reveal a lower bound, since a large planet orbiting at a high angle to the line of sight will perturb its star radially as much as a much smaller planet with an orbital plane on the line of sight. It has been suggested that planets with high eccentricities calculated by this method may in fact be two-planet systems of circular or near-circular resonant orbit; the radial velocity method to detect exoplanets is based on the detection of variations in the velocity of the central star, due to the changing direction of the gravitational pull from an exoplanet as it orbits the star. When the star moves towards us, its spectrum is blueshifted, while it is redshifted when it moves away from us. By looking at the spectrum of a star—and so, measuring its velocity—it can be determined if it moves periodically due to the influence of an exoplanet companion. From the instrumental perspective, velocities are measured relative to the telescope's motion. So an important first step of the data reduction is to remove the contributions of the Earth's elliptic motion around the sun at ± 30 km/s, a monthly rotation of ± 13 m/s of the Earth around the center of gravity of the Earth-Moon system, the daily rotation of the telescope with the Earth crust around the Earth axis, up to ±460 m/s at the equator and proportional to the cosine of the telescope's geographic latitude, small contributions from the Earth polar motion at the level of mm/s, contributions of 230 km/s from the motion around the Galactic center and associated proper motions.
In the case of spectroscopic measurements corrections of the order of ±20 cm/s with respect to aberration. Proper motion Peculiar velocity Relative velocity Space velocity The Radial Velocity Equation in the Search for Exoplanets
Stellar parallax is the apparent shift of position of any nearby star against the background of distant objects. Created by the different orbital positions of Earth, the small observed shift is largest at time intervals of about six months, when Earth arrives at opposite sides of the Sun in its orbit, giving a baseline distance of about two astronomical units between observations; the parallax itself is considered to be half of this maximum, about equivalent to the observational shift that would occur due to the different positions of Earth and the Sun, a baseline of one astronomical unit. Stellar parallax is so difficult to detect that its existence was the subject of much debate in astronomy for hundreds of years, it was first observed in 1806 by Giuseppe Calandrelli who reported parallax in α-Lyrae in his work "Osservazione e riflessione sulla parallasse annua dall’alfa della Lira". In 1838 Friedrich Bessel made the first successful parallax measurement, for the star 61 Cygni, using a Fraunhofer heliometer at Königsberg Observatory.
Once a star's parallax is known, its distance from Earth can be computed trigonometrically. But the more distant an object is, the smaller its parallax. With 21st-century techniques in astrometry, the limits of accurate measurement make distances farther away than about 100 parsecs too approximate to be useful when obtained by this technique; this limits the applicability of parallax as a measurement of distance to objects that are close on a galactic scale. Other techniques, such as spectral red-shift, are required to measure the distance of more remote objects. Stellar parallax measures are given in the tiny units of arcseconds, or in thousandths of arcseconds; the distance unit parsec is defined as the length of the leg of a right triangle adjacent to the angle of one arcsecond at one vertex, where the other leg is 1 AU long. Because stellar parallaxes and distances all involve such skinny right triangles, a convenient trigonometric approximation can be used to convert parallaxes to distance.
The approximate distance is the reciprocal of the parallax: d ≃ 1 / p. For example, Proxima Centauri, whose parallax is 0.7687, is 1 / 0.7687 parsecs = 1.3009 parsecs distant. Stellar parallax is so small that its apparent absence was used as a scientific argument against heliocentrism during the early modern age, it is clear from Euclid's geometry that the effect would be undetectable if the stars were far enough away, but for various reasons such gigantic distances involved seemed implausible: it was one of Tycho Brahe's principal objections to Copernican heliocentrism that in order for it to be compatible with the lack of observable stellar parallax, there would have to be an enormous and unlikely void between the orbit of Saturn and the eighth sphere. James Bradley first tried to measure stellar parallaxes in 1729; the stellar movement proved too insignificant for his telescope, but he instead discovered the aberration of light and the nutation of Earth's axis, catalogued 3222 stars. Stellar parallax is most measured using annual parallax, defined as the difference in position of a star as seen from Earth and Sun, i.e. the angle subtended at a star by the mean radius of Earth's orbit around the Sun.
The parsec is defined as the distance. Annual parallax is measured by observing the position of a star at different times of the year as Earth moves through its orbit. Measurement of annual parallax was the first reliable way to determine the distances to the closest stars; the first successful measurements of stellar parallax were made by Friedrich Bessel in 1838 for the star 61 Cygni using a heliometer. Being difficult to measure, only about 60 stellar parallaxes had been obtained by the end of the 19th century by use of the filar micrometer. Astrographs using astronomical photographic plates sped the process in the early 20th century. Automated plate-measuring machines and more sophisticated computer technology of the 1960s allowed more efficient compilation of star catalogues. In the 1980s, charge-coupled devices replaced photographic plates and reduced optical uncertainties to one milliarcsecond. Stellar parallax remains the standard for calibrating other measurement methods. Accurate calculations of distance based on stellar parallax require a measurement of the distance from Earth to the Sun, now known to exquisite accuracy based on radar reflection off the surfaces of planets.
The angles involved in these calculations are small and thus difficult to measure. The nearest star to the Sun, Proxima Centauri, has a parallax of 0.7687 ± 0.0003 arcsec. This angle is that subtended by an object 2 centimeters in diameter located 5.3 kilometers away. In 1989 the satellite Hipparcos was launched for obtaining parallaxes and proper motions of nearby stars, increasing the number of stellar parallaxes measured to milliarcsecond accuracy a thousandfold. So, Hipparcos is only able to measure parallax angles for stars up to about 1,600 light-years away, a little more than one percent of the diameter of the Milky Way Galaxy; the Hubble telescope WFC3 now has a precision of 20 to 40 microarcseconds, enabling reliable distance measurements u
A variable star is a star whose brightness as seen from Earth fluctuates. This variation may be caused by a change in emitted light or by something blocking the light, so variable stars are classified as either: Intrinsic variables, whose luminosity changes. Extrinsic variables, whose apparent changes in brightness are due to changes in the amount of their light that can reach Earth. Many most, stars have at least some variation in luminosity: the energy output of our Sun, for example, varies by about 0.1% over an 11-year solar cycle. An ancient Egyptian calendar of lucky and unlucky days composed some 3,200 years ago may be the oldest preserved historical document of the discovery of a variable star, the eclipsing binary Algol. Of the modern astronomers, the first variable star was identified in 1638 when Johannes Holwarda noticed that Omicron Ceti pulsated in a cycle taking 11 months; this discovery, combined with supernovae observed in 1572 and 1604, proved that the starry sky was not eternally invariable as Aristotle and other ancient philosophers had taught.
In this way, the discovery of variable stars contributed to the astronomical revolution of the sixteenth and early seventeenth centuries. The second variable star to be described was the eclipsing variable Algol, by Geminiano Montanari in 1669. Chi Cygni was identified in 1686 by G. Kirch R Hydrae in 1704 by G. D. Maraldi. By 1786 ten variable stars were known. John Goodricke himself discovered Beta Lyrae. Since 1850 the number of known variable stars has increased especially after 1890 when it became possible to identify variable stars by means of photography; the latest edition of the General Catalogue of Variable Stars lists more than 46,000 variable stars in the Milky Way, as well as 10,000 in other galaxies, over 10,000'suspected' variables. The most common kinds of variability involve changes in brightness, but other types of variability occur, in particular changes in the spectrum. By combining light curve data with observed spectral changes, astronomers are able to explain why a particular star is variable.
Variable stars are analysed using photometry, spectrophotometry and spectroscopy. Measurements of their changes in brightness can be plotted to produce light curves. For regular variables, the period of variation and its amplitude can be well established. Peak brightnesses in the light curve are known as maxima. Amateur astronomers can do useful scientific study of variable stars by visually comparing the star with other stars within the same telescopic field of view of which the magnitudes are known and constant. By estimating the variable's magnitude and noting the time of observation a visual lightcurve can be constructed; the American Association of Variable Star Observers collects such observations from participants around the world and shares the data with the scientific community. From the light curve the following data are derived: are the brightness variations periodical, irregular, or unique? What is the period of the brightness fluctuations? What is the shape of the light curve? From the spectrum the following data are derived: what kind of star is it: what is its temperature, its luminosity class? is it a single star, or a binary? does the spectrum change with time?
Changes in brightness may depend on the part of the spectrum, observed if the wavelengths of spectral lines are shifted this points to movements strong magnetic fields on the star betray themselves in the spectrum abnormal emission or absorption lines may be indication of a hot stellar atmosphere, or gas clouds surrounding the star. In few cases it is possible to make pictures of a stellar disk; these may show darker spots on its surface. Combining light curves with spectral data gives a clue as to the changes that occur in a variable star. For example, evidence for a pulsating star is found in its shifting spectrum because its surface periodically moves toward and away from us, with the same frequency as its changing brightness. About two-thirds of all variable stars appear to be pulsating. In the 1930s astronomer Arthur Stanley Eddington showed that the mathematical equations that describe the interior of a star may lead to instabilities that cause a star to pulsate; the most common type of instability is related to oscillations in the degree of ionization in outer, convective layers of the star.
Suppose the star is in the swelling phase. Its outer layers expand; because of the decreasing temperature the degree of ionization decreases. This makes the gas more transparent, thus makes it easier for the star to radiate its energy; this in turn will make the star start to contract. As the gas is thereby compressed, it is heated and the degree of ionization again increases. Thi
In astronomy, luminosity is the total amount of energy emitted per unit of time by a star, galaxy, or other astronomical object. As a term for energy emitted per unit time, luminosity is synonymous with power. In SI units luminosity is measured in joules per second or watts. Values for luminosity are given in the terms of the luminosity of the Sun, L⊙. Luminosity can be given in terms of the astronomical magnitude system: the absolute bolometric magnitude of an object is a logarithmic measure of its total energy emission rate, while absolute magnitude is a logarithmic measure of the luminosity within some specific wavelength range or filter band. In contrast, the term brightness in astronomy is used to refer to an object's apparent brightness: that is, how bright an object appears to an observer. Apparent brightness depends on both the luminosity of the object and the distance between the object and observer, on any absorption of light along the path from object to observer. Apparent magnitude is a logarithmic measure of apparent brightness.
The distance determined by luminosity measures can be somewhat ambiguous, is thus sometimes called the luminosity distance. In astronomy, luminosity is the amount of electromagnetic energy; when not qualified, the term "luminosity" means bolometric luminosity, measured either in the SI units, watts, or in terms of solar luminosities. A bolometer is the instrument used to measure radiant energy over a wide band by absorption and measurement of heating. A star radiates neutrinos, which carry off some energy, contributing to the star's total luminosity; the IAU has defined a nominal solar luminosity of 3.828×1026 W to promote publication of consistent and comparable values in units of the solar luminosity. While bolometers do exist, they cannot be used to measure the apparent brightness of a star because they are insufficiently sensitive across the electromagnetic spectrum and because most wavelengths do not reach the surface of the Earth. In practice bolometric magnitudes are measured by taking measurements at certain wavelengths and constructing a model of the total spectrum, most to match those measurements.
In some cases, the process of estimation is extreme, with luminosities being calculated when less than 1% of the energy output is observed, for example with a hot Wolf-Rayet star observed only in the infra-red. Bolometric luminosities can be calculated using a bolometric correction to a luminosity in a particular passband; the term luminosity is used in relation to particular passbands such as a visual luminosity of K-band luminosity. These are not luminosities in the strict sense of an absolute measure of radiated power, but absolute magnitudes defined for a given filter in a photometric system. Several different photometric systems exist; some such as the UBV or Johnson system are defined against photometric standard stars, while others such as the AB system are defined in terms of a spectral flux density. A star's luminosity can be determined from two stellar characteristics: size and effective temperature; the former is represented in terms of solar radii, R⊙, while the latter is represented in kelvins, but in most cases neither can be measured directly.
To determine a star's radius, two other metrics are needed: the star's angular diameter and its distance from Earth. Both can be measured with great accuracy in certain cases, with cool supergiants having large angular diameters, some cool evolved stars having masers in their atmospheres that can be used to measure the parallax using VLBI. However, for most stars the angular diameter or parallax, or both, are far below our ability to measure with any certainty. Since the effective temperature is a number that represents the temperature of a black body that would reproduce the luminosity, it cannot be measured directly, but it can be estimated from the spectrum. An alternative way to measure stellar luminosity is to measure the star's apparent brightness and distance. A third component needed to derive the luminosity is the degree of interstellar extinction, present, a condition that arises because of gas and dust present in the interstellar medium, the Earth's atmosphere, circumstellar matter.
One of astronomy's central challenges in determining a star's luminosity is to derive accurate measurements for each of these components, without which an accurate luminosity figure remains elusive. Extinction can only be measured directly if the actual and observed luminosities are both known, but it can be estimated from the observed colour of a star, using models of the expected level of reddening from the interstellar medium. In the current system of stellar classification, stars are grouped according to temperature, with the massive young and energetic Class O stars boasting temperatures in excess of 30,000 K while the less massive older Class M stars exhibit temperatures less than 3,500 K; because luminosity is proportional to temperature to the fourth power, the large variation in stellar temperatures produces an vaster variation in stellar luminosity. Because the luminosity depends on a high power of the stellar mass, high mass luminous stars have much shorter lifetimes; the most luminous stars are always young stars, no more than a few million years for the most extreme.
In the Hertzsprung–Russell diagram, the x-axis represents temperature or spectral type while the y-axis represents luminosity or magnitude. The vast majority of stars are found along the main sequence with blue Class O stars found at the top left of the chart while red Class M stars fall to the bottom right. Certain stars like Deneb and Betelgeuse are
ArXiv is a repository of electronic preprints approved for posting after moderation, but not full peer review. It consists of scientific papers in the fields of mathematics, astronomy, electrical engineering, computer science, quantitative biology, mathematical finance and economics, which can be accessed online. In many fields of mathematics and physics all scientific papers are self-archived on the arXiv repository. Begun on August 14, 1991, arXiv.org passed the half-million-article milestone on October 3, 2008, had hit a million by the end of 2014. By October 2016 the submission rate had grown to more than 10,000 per month. ArXiv was made possible by the compact TeX file format, which allowed scientific papers to be transmitted over the Internet and rendered client-side. Around 1990, Joanne Cohn began emailing physics preprints to colleagues as TeX files, but the number of papers being sent soon filled mailboxes to capacity. Paul Ginsparg recognized the need for central storage, in August 1991 he created a central repository mailbox stored at the Los Alamos National Laboratory which could be accessed from any computer.
Additional modes of access were soon added: FTP in 1991, Gopher in 1992, the World Wide Web in 1993. The term e-print was adopted to describe the articles, it began as a physics archive, called the LANL preprint archive, but soon expanded to include astronomy, computer science, quantitative biology and, most statistics. Its original domain name was xxx.lanl.gov. Due to LANL's lack of interest in the expanding technology, in 2001 Ginsparg changed institutions to Cornell University and changed the name of the repository to arXiv.org. It is now hosted principally with eight mirrors around the world, its existence was one of the precipitating factors that led to the current movement in scientific publishing known as open access. Mathematicians and scientists upload their papers to arXiv.org for worldwide access and sometimes for reviews before they are published in peer-reviewed journals. Ginsparg was awarded a MacArthur Fellowship in 2002 for his establishment of arXiv; the annual budget for arXiv is $826,000 for 2013 to 2017, funded jointly by Cornell University Library, the Simons Foundation and annual fee income from member institutions.
This model arose in 2010, when Cornell sought to broaden the financial funding of the project by asking institutions to make annual voluntary contributions based on the amount of download usage by each institution. Each member institution pledges a five-year funding commitment to support arXiv. Based on institutional usage ranking, the annual fees are set in four tiers from $1,000 to $4,400. Cornell's goal is to raise at least $504,000 per year through membership fees generated by 220 institutions. In September 2011, Cornell University Library took overall administrative and financial responsibility for arXiv's operation and development. Ginsparg was quoted in the Chronicle of Higher Education as saying it "was supposed to be a three-hour tour, not a life sentence". However, Ginsparg remains on the arXiv Scientific Advisory Board and on the arXiv Physics Advisory Committee. Although arXiv is not peer reviewed, a collection of moderators for each area review the submissions; the lists of moderators for many sections of arXiv are publicly available, but moderators for most of the physics sections remain unlisted.
Additionally, an "endorsement" system was introduced in 2004 as part of an effort to ensure content is relevant and of interest to current research in the specified disciplines. Under the system, for categories that use it, an author must be endorsed by an established arXiv author before being allowed to submit papers to those categories. Endorsers are not asked to review the paper for errors, but to check whether the paper is appropriate for the intended subject area. New authors from recognized academic institutions receive automatic endorsement, which in practice means that they do not need to deal with the endorsement system at all. However, the endorsement system has attracted criticism for restricting scientific inquiry. A majority of the e-prints are submitted to journals for publication, but some work, including some influential papers, remain purely as e-prints and are never published in a peer-reviewed journal. A well-known example of the latter is an outline of a proof of Thurston's geometrization conjecture, including the Poincaré conjecture as a particular case, uploaded by Grigori Perelman in November 2002.
Perelman appears content to forgo the traditional peer-reviewed journal process, stating: "If anybody is interested in my way of solving the problem, it's all there – let them go and read about it". Despite this non-traditional method of publication, other mathematicians recognized this work by offering the Fields Medal and Clay Mathematics Millennium Prizes to Perelman, both of which he refused. Papers can be submitted in any of several formats, including LaTeX, PDF printed from a word processor other than TeX or LaTeX; the submission is rejected by the arXiv software if generating the final PDF file fails, if any image file is too large, or if the total size of the submission is too large. ArXiv now allows one to store and modify an incomplete submission, only finalize the submission when ready; the time stamp on the article is set. The standard access route is through one of several mirrors. Sev
Auriga is one of the 88 modern constellations. Located north of the celestial equator, its name is the Latin word for “the charioteer”, associating it with various mythological beings, including Erichthonius and Myrtilus. Auriga is most prominent during winter evenings in the northern Hemisphere, along with the five other constellations that have stars in the Winter Hexagon asterism; because of its northern declination, Auriga is only visible in its entirety as far as 34° south. A large constellation, with an area of 657 square degrees, it is half the size of the largest constellation, Hydra, its brightest star, Capella, is an unusual multiple star system among the brightest stars in the night sky. Beta Aurigae is an interesting variable star in the constellation; because of its position near the winter Milky Way, Auriga has many bright open clusters in its borders, including M36, M37, M38, popular targets for amateur astronomers. In addition, it has one prominent nebula, the Flaming Star Nebula, associated with the variable star AE Aurigae.
In Chinese mythology, Auriga's stars were incorporated into several constellations, including the celestial emperors' chariots, made up of the modern constellation's brightest stars. Auriga is home to the radiant for the Aurigids, Zeta Aurigids, Delta Aurigids, the hypothesized Iota Aurigids; the first record of Auriga's stars was in Mesopotamia as a constellation called GAM, representing a scimitar or crook. However, this may have represented just the modern constellation as a whole. GAM in the MUL. APIN; the crook of Auriga shepherd. It was formed from most of the stars of the modern constellation. Bedouin astronomers created constellations that were groups of animals, where each star represented one animal; the stars of Auriga comprised a herd of goats, an association present in Greek mythology. The association with goats carried into the Greek astronomical tradition, though it became associated with a charioteer along with the shepherd. In Greek mythology, Auriga is identified as the mythological Greek hero Erichthonius of Athens, the chthonic son of Hephaestus, raised by the goddess Athena.
Erichthonius was credited to be the inventor of the quadriga, the four-horse chariot, which he used in the battle against the usurper Amphictyon, the event that made Erichthonius the king of Athens. His chariot was created in the image of the Sun's chariot, the reason Zeus placed him in the heavens; the Athenian hero dedicated himself to Athena and, soon after, Zeus raised him into the night sky in honor of his ingenuity and heroic deeds. Auriga, however, is sometimes described as Myrtilus, Hermes's son and the charioteer of Oenomaus; the association of Auriga and Myrtilus is supported by depictions of the constellation, which show a chariot. Myrtilus's chariot was destroyed in a race intended for suitors to win the heart of Oenomaus's daughter Hippodamia. Myrtilus earned his position in the sky when Hippodamia's successful suitor, killed him, despite his complicity in helping Pelops win her hand. After his death, Myrtilus's father Hermes placed him in the sky, yet another mythological association of Auriga is Theseus's son Hippolytus.
He was ejected from Athens after he refused the romantic advances of his stepmother Phaedra, who committed suicide as a result. He was revived by Asclepius. Regardless of Auriga's specific representation, it is that the constellation was created by the ancient Greeks to commemorate the importance of the chariot in their society. An incidental appearance of Auriga in Greek mythology is as the limbs of Medea's brother. In the myth of Jason and the Argonauts, as they journeyed home, Medea killed her brother and dismembered him, flinging the parts of his body into the sea, represented by the Milky Way; each individual star represents a different limb. Capella is associated with the mythological she-goat Amalthea, it forms an asterism with the stars Epsilon Aurigae, Zeta Aurigae, Eta Aurigae, the latter two of which are known as the Haedi. Though most associated with Amalthea, Capella has sometimes been associated with Amalthea's owner, a nymph; the myth of the nymph says that the goat's hideous appearance, resembling a Gorgon, was responsible for the Titans' defeat, because Zeus skinned the goat and wore it as his aegis.
The asterism containing the three goats had been a separate constellation. Before that, Capella was sometimes seen as its own constellation—by Pliny the Elder and Manilius—called Capra, Caper, or Hircus, all of which relate to its status as the "goat star". Zeta Aurigae and Eta Aurigae were first called the "Kids" by Cleostratus, an ancient Greek astronomer. Traditionally, illustrations of Auriga represent it as its driver; the charioteer has two kids under his left arm. However, depictions of Auriga have been inconsistent over the years; the reins in his right hand have been drawn as a whip, though Capella is always over his left shoulder and the Kids under his left arm. The 1488 atlas Hyginus deviated from this typical depiction by showing a four-wheeled cart driven by Auriga
A shell star is a star having a spectrum that shows broad absorption lines, plus some narrow absorption lines. They also show some emission lines from the Balmer series but of other lines; the broad absorption lines are due to rapid rotation of the photosphere, the emission lines from an equatorial disk, the narrow absorption lines are produced when the disc is seen nearly edge-on. Shell stars have spectral types O7.5 to F5, with rotation velocities of 200–300 km/s, not far from the point when the rotational acceleration would disrupt the star. The shell stars are defined as a group by the existence of rotationally broadened photospheric spectral lines in combination with narrow absorption lines. Emission lines are present but not regarded as a defining feature; the exact spectral lines present vary to some extent: Balmer emission lines are common, but may be weak or absent in cooler stars. The photospheric lines are rotationally broadened showing projected velocities of 200 km/s or more; the line profiles in shell star spectra are complex, with variable wings and superpositions of absorption and emission features.
In some cases, particular absorption of emission features are only visible as modifications to a line profile, or a weakening of another line. This leads to asymmetric lines. Shell stars commonly have emission lines and thus are Be stars, although they can occur in spectral classes O, A, F. Shell stars have been subdivided into four categories, although these categories are no longer recognised as meaningful and are seen in modern publications: early Be stars of spectral types O7.5 to B2.5 middle Be stars of types B3 to B6.5, late Be stars of types B7 to B9.5, A-F shell stars ranging from A0 to F5. The vast majority of known shell stars are of spectral class B; however because of this many cooler shell stars have remained undetected. The Be phenomenon, hence the term Be star itself, is now applied to similar stars not only of spectral class B, but A and sometimes O and F. Shell stars show variability in their spectra and brightness; the shell features may come and go, with the star changing from a shell star to a normal B star or Be star.
Shell stars which show irregular variability due to changes in, or the disappearance of, the "shell" are called Gamma Cassiopeiae variables. Pleione and Gamma Cassiopeiae itself are both variable stars that have intermittent shell episodes where strong shell features appear in the spectrum and the brightness increases or decreases significantly. At other times the shell is not detectable in the spectrum, the emission lines may disappear. Porter, John M.. "On the rotational velocities of Be and Be-shell stars". Monthly Notices of the Royal Astronomical Society. 280: L31–L35. Bibcode:1996MNRAS.280L..31P. Doi:10.1093/mnras/280.3. L31. Quirrenbach, A.. "Maximum-entropy maps of the Be shell star zeta Tauri from optical long-baseline interferometry". Astronomy and Astrophysics. 283: L13. Bibcode:1994A&A...283L..13Q