The variability of stars existing within our universe represents diversity and uniqueness. Humanity may not observe the whole life of a single star, however, techniques have been proposed to study the stars. We have different ways how to measure a star's temperature, age, and mass even, even though it seems impossible.
The said variability of stars is key to the measurement of a star's characteristics. The existence of variable stars suggests that stars change on a time interval, unfortunately, too long humanity cannot observe. Nevertheless, humanity have developed a way of understanding the stars through their variability. Needless to say, variability provides unique physical properties every time we study specific stars, and their variations have been the "experimental laboratories" for stellar physics (Templeton, 2010).
Before we can hypothesize about the evolution of stars, it is necessary to measure the mass and the age of the stars. These quantities are deemed fundamental in understanding the stars because if we can measure the age and mass of variable stars, then we can provide more evidence significant to hypothesize about the evolution of stars. The mass represents the energy of the stars while age represents the status of the stars in their evolution. However, these two quantities are difficult to measure directly. Fortunately, two characteristics are easily observed: (1) color and (2) brightness (Young, Freedman & Ford, 2012).
During the early 20th century, two astronomers have discovered patterns with the color and the brightness of variable stars. Ejnar Hertzsprung and Henry Norris Russell have found out that the plot of the brightness of the variable stars versus their spectral color resembles well-defined areas. This means that a particular star with a certain brightness could only have a certain range of colors while a particular star with a certain spectral color can only have a certain range of brightness. What is more curious about the Hertzsprung-Russell diagram is that the plot represents an image of the evolutionary states of the stars. The well-defined areas within the Hertzsprung-Russell diagram also represent the physical properties of the stars in general. By classifying and joining all stages of stars into one picture provides a broader image of stellar evolution (de Boer & Seggewis, 2008).
The birth of stars occurs on star forming regions of concentrations of gas and dust, called nebulae, collapsing under the gravity of forming stars. These young variable stars, called protostars, are undergoing a stage where accretion of dust and gas from the nebulae primarily occurs. As they accumulate mass and energy, they enter a stage wherein they are still weakly accreting materials from the nebulae. However, they are less bright relative to stars of their size and color, are near gaseous nebulae, and emit light from highly excited atoms of a thin gas. This stage is often named the pre-main sequence or PMS stars (Schulz, 2005).
The characteristics of the PMS stars are due primarily to the accretion. As gravitational potential energy converts to kinetic energy, the gas also gains temperature to the friction. The temperature increases as the gas travel towards the surface, and as it does, it gives off a light. The energy source of pre-main sequence stars results from the proton-proton chain reaction. These nuclear fusion reactions result from the interaction of the gravitational potential energy and the very high temperature especially within the star. The temperature could reach 10 million Kelvins at the core while 4500 Kelvins at the surface (Smith, 1999).
As the star continues to balance the nuclear fusion inside the core and the gravitational energy, the star enters the main sequence stage, the longest stage in its evolution. Within the Hertzsprung-Russell diagram, stars in the main sequence stage lie on a band starting from the lower boundary named the Zero-Age main sequence. The regulation of the energy production of a main-sequence star is due to physical principles: (1) thermal radiation law, which determines the resultant energy required for nuclear energy production, (2) reaction rate for the proton-proton chain reaction, and (3) the balance of inward pull of gravity by the pressure determined by Ideal Gas Law (Templeton, 2010). Therefore, the main-sequence stars are in an equilibrium wherein the energy is thoroughly balanced and regulated. For instance, an insufficient energy in the core means that the star has to adjust from the star's gravitational energy. The core would contract, heat up, build up pressure, and then, the generation rate of nuclear energy would speed up until it matches the radiating energy. The opposite happens when an excessive energy in the core happens to mismatch the energy radiating from the surface (Shaviv, 2009).
It is indeed the longest stage during the evolution of a star. However, it is important to note that the entire evolution of the star depends on its mass. The mass of the stars represents the total energy it can produce and the rate it can produce the energy (Lada & Kylafis, 1991).
Eventually, the hydrogen in the core will be used up due to the nuclear fusion. The result will an enormous amount of Helium. As the nuclear reaction slows down, the gravity will take over. The core shrinks while the layers outside the core collapse. The compression from the collapse of the layers, called shell layer, heats up the gases enough for fusion to start. The fusion becomes rapid because of the temperature increase due to the continuous compression. The luminosity increases as the size of the star increases. The gas surrounding the core are being pushed away due to the outward pressure, and the star expands until it becomes a red giant (Shaviv, 2009).
However, this is not always the case. Depending on the mass of the star during its main-sequence stage, the star would end the main-sequence stage into three possible scenarios. For the case of ordinary lower mass stars, including our sun, the end of the main-sequence stage means that it would leave a white dwarf star after the red giant phase. As the star exhausts its helium core, the energy would not suffice the pressure of the expansion, then, it will start to lose mass from the powerful stellar winds. As the star loses its mass, a stage called the post-asymptomatic giant branch, it will leave behind its hot carbon core surrounded by nebular gas. The carbon core will cool, and will develop into a white dwarf star (Delehanty, nd). As the name suggests, the white dwarf stars are smaller remains of the red giant phase. Typically, their diameters are just thousands of kilometers wide. However, the reason they shine is not due to nuclear fusion but from the leftover heat from when it was in its main-sequence stage. The stored thermal energy would die off eventually and become a black dwarf (Smith, 1999).
As for heavier stars than our sun, massive stars perish more dramatically than less dense stars. Although massive stars also exhaust their helium core, the nuclear fusion still proceeds as elements such as carbon transform into heavier elements. As the nuclear fusion cycle reaches its end in the form of iron, the collapse of the star cannot be balanced by any source of energy. The compression of the core reaches nuclear densities so high that it resists further collapse, leaving all in-falling elements to “bounce off”. This “bouncing off” of materials together with the release of energetic neutrinos become observable as a supernova explosion. This supernova explosion releases heavier elements into interstellar space, leaving a massive volume of elements into vacuüm. Typically, this is where other stars start to form from the leftovers of the supernova explosion. However, depending on the mass of the progenitor star, the core could become either a neutron star or a black hole (Templeton, 2010). A neutron star is detectable due to its high radio emission as pulsars.
Figure 1. The size of the star suggests the outcome of its evolution. Source: Dunbar, 2009
Still, the resulting nebulae from the death of the stars become the birth grounds of the succeeding stars. One interesting idea is that when the death of a star results into a white dwarf, the released materials become ingredients for another star. The result is then a binary system of a white dwarf and another star. For instance, the bright star Sirius has a white dwarf companion (Smith, 1999). Still, regardless of the outcome of the star, the cycle continues.
The variability of the stars become essential to studying the stars. The different physical characteristics provide unique evidence suggesting that stars always change in its lifespan. The observable characteristics of the stars, their color, and brightness, becomes fundamental in concocting hypotheses concerning the stellar evolution.
The birth of the stars starts in a massive cloud of dust and gas. As a protostar evolves, it starts to accumulate mass by accreting the surrounding materials. As it gains mass, the pressure and temperature starts to increase until it is sufficient enough to start proton-proton chain reaction. The nuclear fusion becomes stable but is still weak in its PMS star stage. It continues to accumulate materials until it enters its zero-age main-sequence stage. The nuclear fusion reaction becomes stable, and an equilibrium occurs between gravitational energy and nuclear energy. However, the resources are limited, and eventually, they will be exhausted. Depending on the mass of the star, the result from the red giant phase could be a white dwarf, a neutron star, and a black hole. Nonetheless, the death of the star provides the materials sufficient for the birth of another star.
de Boer, K.S. & Seggewiss, W. (2008). Stars and Stellar Evolution. Les Ulis Cedex A, France: EDP Sciences.
Delehanty, Mark. (n.d.). "Stellar Evolution, the lives of stars." AstronomyToday. Retrieved from http://www.astronomytoday.com/cosmology/evol.html
Dunbar, Brian. (10 April 2009). "Stellar Evolution The Birth, Life, and Death of a Star." NASA. Retrieved from http://www.nasa.gov/audience/forstudents/912/features/stellar_evol_feat_912.html
Lada, C.J. & Kylafis, N.D. (1991). The Physics of Star Formation and Early Stellar Evolution. Berlin, Germany: Springer Science + Business Media Dordrecht.
Redd, N.T. (3 April 2015). "Hot, Young Star 'Missing Link' of Stellar Evolution." SPACE.com. Retrieved from http://www.space.com/29014massivestarevolutionmissinglink.html
Schulz, Norbert S. (2005). From Dust to Stars: Studies of the Formation and Early Evolution of Stars. Chichester, UK: Praxis Publishing Ltd.
Shaviv, Giora. (2009). The Life of Stars: The Controversial Inception and Emergence of the Theory of Stellar Structure. Berlin, Germany: Springer-Verlag.
Smith, H.E. (16 April 1999). "Stellar Evolution I - Solar Type Stars." University of California, San Diego Center for Astrophysics & Space Sciences. Retrieved from http://casswww.ucsd.edu/archive/public/tutorial/StevI.html
Templeton, Matthew. (13 September 2010). "Stellar Evolution." Retrieved from http://www.aavso.org/stellarevolution
Young, H.D., Freedman, R.A., & Ford, A.L. (2012). Sears and Zemansky's University Physics, 13th Edition. San Francisco, CA: Pearson Education, Inc., publishing as Addison-Wesley.