We finally made it! Our example T Tauri star has now contracted enough to sustain sufficient temperatures to fuse hydrogen, thereby classifying it as a main sequence star. Stars remain on the sequence for the grand majority of their lifetime. The specific amount of time spent on the main sequence, called the main sequence lifetime, can be calculated and depends completely on the mass of the star at hand. Stars with lower mass stay in the main sequence for significantly longer (on the order of billions of years longer) than those with lots of mass. For example, a star with high mass may stay on the main sequence for a short period of 100 million years (if very very massive) whereas our sun, with a low mass, will stay in the main sequence for 10 billion years (5 billion more years from today). In this coming section, we will go in depth into the specific properties of stars in main sequence so you can be an expert for your next test!
Stellar Fusion, Lifetime, & Progression
It is important to first note where the main sequence on the HR diagram. The main sequence is a upward sloping line in the center of the HR diagram, expressing a positive relationship between temperature and luminosity (note that as you move left on a HR diagram, temperatures get higher). Thus, stars higher on the main sequence are capable of sustaining higher temperatures and thus are stars of higher mass. Once stars land on their zero-age main-sequence point from either the Hayashi or Henyey tract, they DO NOT move up or down the main sequence line (common misconception). They stay at their one point until they begin to run our of their hydrogen fuel and evolve into another form (subgiant for low mass stars, luminous blue giant or supergiant for high mass stars because their subgiant stage is incredibly short or non-existent).
A question you may have is why, or even how can a star remain so fixed on the main sequence for such a long period of time? The answer is simple: it has a lot of fuel. And I mean a lot of it. Our sun alone, which is considered "low mass" has roughly 1.99*10^30 kilograms of hydrogen fuel to burn for the rest of its lifetime. For reference, that is over 1000 times the mass of all our the planets, comets, and asteroids in our solar system combined (2.66*10^27 kg). In order to use this crazy amount of fuel efficiently, stars undergo a variety of different processes throughout its evolution.
Proton-Proton-Chain
The proton-proton chain is the dominant nuclear fusion process in stars with masses similar to or less than that of the Sun, typically with core temperatures below about 15 million Kelvin. This chain reaction is the primary energy source in these stars during their time on the main sequence, where they spend the majority of their lives. While more massive stars rely primarily on the carbon-nitrogen-oxygen (CNO) cycle, which requires the presence of heavier nuclei and higher temperatures, smaller stars lack the core conditions to sustain such reactions efficiently. Instead, they convert hydrogen into helium through a sequence of reactions involving only protons and light isotopes. The process begins at the extreme pressures and temperatures found in the stellar core, where protons—hydrogen nuclei—have enough kinetic energy to overcome the electrostatic repulsion between them. When they do fuse, the resulting nuclear transformations release vast amounts of energy that heat the star and provide the outward pressure needed to counteract the force of gravity. The balance between gravitational collapse and radiation pressure, known as hydrostatic equilibrium, is maintained through the energy generated by the proton-proton chain.
The first step in the proton-proton chain is both rare and incredibly important, as it determines the overall pace of energy generation in the star. Two protons collide and fuse to form a deuteron, a nucleus consisting of one proton and one neutron. However, because this process involves the weak nuclear force—a proton must transform into a neutron via beta-plus decay—it has a very low probability of occurring. This rarity is beneficial in terms of stellar lifetimes; if it occurred more frequently, stars like the Sun would burn through their hydrogen fuel much more quickly. The full reaction is:
\[^1_1\text{H} + ^1_1\text{H} \rightarrow\ ^2_1\text{D} + e^+ + \nu_e\]
Here, a positron (\(e^+\)) and an electron neutrino (\(\nu_e\)) are emitted. The positron rapidly annihilates with a nearby electron, producing two gamma-ray photons and contributing additional energy to the system. The neutrino, on the other hand, barely interacts with matter and escapes the star entirely, carrying away about 0.26 MeV of the total \(\sim 1.44\) MeV energy released in this step. Although it accounts for only a small percentage of the star’s total energy output, this energy loss to neutrinos has significant implications for stellar energy balance and neutrino detection in solar observatories on Earth.
Once deuterium is formed, it quickly undergoes fusion with another proton to form helium-3. This reaction is much faster than the first, because it does not involve the weak interaction; instead, it proceeds via the strong nuclear force. The reaction is as follows: \[^2_1\text{D} + ^1_1\text{H} \rightarrow\ ^3_2\text{He} + \gamma\] This step releases about 5.49 MeV of energy, primarily in the form of a high-energy gamma-ray photon. These photons interact with surrounding plasma and are gradually degraded in energy as they migrate outward through the radiative zone of the star. It can take thousands to millions of years for these photons to reach the surface and escape as visible light. The \(^3_2\text{He}\) nucleus now contains two protons and one neutron, and will participate in the final step of the chain. It’s important to note that by this point, two protons have already been converted and bound into a light helium isotope, but the process isn’t finished until a stable helium-4 nucleus is formed.
In the most common path of the chain, called the proton-proton I (pp I) chain, two helium-3 nuclei collide to form a helium-4 nucleus, releasing two protons in the process. This reaction is: \[^3_2\text{He} + ^3_2\text{He} \rightarrow\ ^4_2\text{He} + 2^1_1\text{H}\] This step releases approximately 12.86 MeV of energy. The helium-4 nucleus (two protons and two neutrons) is exceptionally stable due to its high binding energy, which is one of the main reasons fusion proceeds toward this product. The two protons that are released can return to the cycle, potentially fusing again in earlier steps. While other branches of the proton-proton chain do exist (namely pp II and pp III), they require higher core temperatures and involve intermediate elements like beryllium-7 and boron-8. These branches are responsible for producing higher-energy neutrinos, such as those detected by experiments like the Sudbury Neutrino Observatory. However, in the Sun, the pp I chain accounts for over 85% of the total fusion reactions.
The net effect of the entire proton-proton chain is the transformation of four protons (hydrogen nuclei) into one helium-4 nucleus, with the emission of two positrons, two neutrinos, and a set of gamma-ray photons. The complete reaction can be summarized as:
\[4^1_1\text{H} \rightarrow\ ^4_2\text{He} + 2e^+ + 2\nu_e + 2\gamma + \text{energy (26.7 MeV)}\]
The total energy released per cycle is about 26.7 MeV, though not all of this is available to heat the star; approximately 2% of it is carried away by neutrinos. The remainder contributes to the thermal energy in the stellar interior, which then radiates outward. Over time, this process gradually converts the hydrogen in the star’s core into helium, altering the core’s composition and energy output. As the hydrogen supply diminishes, the star evolves off the main sequence and begins new fusion processes, but for the vast majority of a star’s life, particularly in stars like the Sun, the proton-proton chain is the foundational process that fuels the star and determines its structure and lifespan.
CNO Cycle
The CNO cycle (carbon-nitrogen-oxygen cycle) is an alternative hydrogen fusion mechanism that dominates in stars more massive than the Sun, typically those with core temperatures exceeding about 17 million Kelvin. While the proton-proton chain is more temperature-efficient at lower energies, the CNO cycle becomes exponentially more efficient as temperature increases, making it the primary energy source for massive main-sequence stars. Unlike the proton-proton chain, which fuses protons directly, the CNO cycle relies on carbon, nitrogen, and oxygen nuclei as catalysts to convert four protons into one helium-4 nucleus. These heavier nuclei are not consumed but serve to facilitate the sequence of reactions. This process illustrates the role of metallicity in stellar evolution; without an initial abundance of carbon, nitrogen, and oxygen—products of previous generations of stars—the CNO cycle cannot function. Thus, the CNO cycle represents a more advanced fusion pathway that only operates efficiently in hotter, more evolved stellar environments.
The CNO cycle is a loop of nuclear reactions that begins and ends with a carbon-12 nucleus. In the first step, a proton fuses with carbon-12 to form nitrogen-13, releasing a gamma-ray photon:
\[^{12}_6\text{C} + ^1_1\text{H} \rightarrow\ ^{13}_7\text{N} + \gamma\]
Nitrogen-13 is an unstable isotope and undergoes beta-plus decay, converting a proton into a neutron and releasing a positron and an electron neutrino:
\[^{13}_7\text{N} \rightarrow\ ^{13}_6\text{C} + e^+ + \nu_e\]
This decay has a half-life of about 10 minutes and produces a carbon-13 nucleus. The positron will annihilate with an electron, emitting gamma-ray photons and contributing to the star’s energy output, while the neutrino escapes. At this point, the cycle has involved the weak interaction (in the beta decay) and released energy from both photon emission and particle annihilation. The importance of weak interaction here parallels the rate-limiting role it plays in the proton-proton chain, determining the overall cycle’s timescale.
Following this, the carbon-13 nucleus captures another proton to form nitrogen-14:
\[^{13}_6\text{C} + ^1_1\text{H} \rightarrow\ ^{14}_7\text{N} + \gamma\]
This is a straightforward strong-force fusion reaction that emits another gamma photon. Nitrogen-14, a stable and abundant isotope, then captures yet another proton to form oxygen-15:
\[^{14}_7\text{N} + ^1_1\text{H} \rightarrow\ ^{15}_8\text{O} + \gamma\]
Oxygen-15 is radioactive and undergoes beta-plus decay with a half-life of about 2 minutes:
\[^{15}_8\text{O} \rightarrow\ ^{15}_7\text{N} + e^+ + \nu_e\]
This forms nitrogen-15, another intermediate in the cycle. Once again, the positron annihilates with an electron to generate gamma radiation, and the neutrino escapes the star, carrying away a small fraction of energy. The cumulative effect of these decays and annihilations is a substantial release of energy, though slightly less than in the proton-proton chain due to different branching ratios and decay rates. The gamma radiation continues to random walk outward through the radiative zone, eventually contributing to the star’s overall luminosity.
In the final step of the main CNO-I cycle, nitrogen-15 fuses with a proton and splits into a helium-4 nucleus and a carbon-12 nucleus:
\[^{15}_7\text{N} + ^1_1\text{H} \rightarrow\ ^{12}_6\text{C} + ^4_2\text{He}\]
This final reaction regenerates the carbon-12 nucleus used in the first step, thus completing the catalytic loop. The net effect of the CNO cycle is the fusion of four protons into one helium-4 nucleus, with the release of two positrons, two neutrinos, and several gamma-ray photons—identical in products to the proton-proton chain, but with a different mechanism and energy distribution. The total energy released in the cycle is approximately 26.7 MeV, similar to the pp-chain, although the temperature sensitivity of the CNO cycle is much steeper: its energy production rate scales roughly with the 17th power of temperature (\(\epsilon \propto T^{17}\)), compared to the fourth power for the pp-chain. As a result, small increases in core temperature in massive stars lead to sharp increases in luminosity.
The CNO cycle has several variants beyond the main loop (CNO-I), including the CNO-II, CNO-III, and CNO-IV cycles, which involve isotopes such as fluorine-17 or neon, and occur at even higher temperatures. These extensions are typically only significant in stars with core temperatures above \(20\) million K. In all these cases, the principle remains the same: hydrogen is fused into helium using heavier nuclei as intermediates or catalysts. The dominance of the CNO cycle in high-mass stars has important implications for stellar structure. For instance, because the reaction rate is so temperature-sensitive, the energy generation in CNO-dominated stars is concentrated very tightly at the core. This leads to steep temperature gradients and the onset of convective cores in more massive stars, unlike the radiative cores of lower-mass, pp-chain dominated stars. The CNO cycle thus not only shapes the way a star shines, but also fundamentally determines its internal structure and subsequent evolutionary pathway. Understanding the CNO cycle is crucial for modeling the lifetimes, luminosities, and nucleosynthetic contributions of massive stars throughout the universe.
Stellar Structure
The specific fusion mechanisms and new higher temperatures of main sequence stars in comparison with their protostar origins allows them to maintain a complex, layered structure. We will now go deep into its nuances so you can be ready for your next Scioly Astro test.
Layers
Corona
The corona is the outermost layer of a star’s atmosphere, extending millions of kilometers into space. It is characterized by its extremely high temperatures, ranging from one to three million Kelvin, which is counterintuitive given its distance from the cooler photosphere below. The source of the corona’s intense heat remains a subject of active research, though it is widely believed to be linked to magnetic reconnection events and wave heating from lower layers of the solar atmosphere. The corona emits strongly in X-rays and ultraviolet radiation, and it is visible from Earth only during a total solar eclipse or with the aid of special instruments like coronagraphs.
Despite its high temperature, the corona has a very low density, making it far less luminous than the photosphere. This low density means that, while individual particles are highly energetic, the total energy output is relatively small. The corona is also the origin of the solar wind, a continuous stream of charged particles that flows outward through the solar system. This wind interacts with planetary magnetospheres, causing auroras and influencing space weather. The corona is structured by the Sun’s magnetic field, forming loops, streamers, and plumes that can be observed in ultraviolet and X-ray wavelengths.
Coronal activity is highly dynamic, with constant changes and eruptions linked to magnetic field interactions. Coronal mass ejections (CMEs) are massive bursts of solar plasma and magnetic fields released into space from the corona, capable of disturbing Earth’s magnetosphere and causing geomagnetic storms. The corona’s magnetic topology also plays a critical role in confining plasma and dictating the shape of solar flares. These phenomena make the corona not just a peripheral feature, but a central player in solar and stellar physics.
Transition Region
The transition region is a narrow, highly stratified layer between the chromosphere and the corona, typically only a few hundred kilometers thick. Despite its limited spatial extent, it is critically important in stellar atmospheric models because it marks a rapid temperature rise from about 20,000 K in the chromosphere to over 1,000,000 K in the corona. This steep gradient indicates a fundamental change in the heating mechanisms at play and suggests intense energy transfer processes, possibly involving magnetohydrodynamic waves and turbulent dissipation.
The transition region is where the gas changes from being largely neutral in the chromosphere to fully ionized in the corona. This ionization leads to a dramatic increase in thermal conductivity, particularly along magnetic field lines, which may explain the abrupt temperature rise. The region emits strongly in ultraviolet lines of ions such as C IV, O V, and Si IV, making it observable with ultraviolet telescopes and spectrometers. Understanding the structure and dynamics of the transition region is essential for connecting chromospheric activity to coronal phenomena.
Because it is so thin, the transition region is highly dynamic and subject to rapid changes driven by underlying magnetic activity. It acts as a bottleneck for mass and energy flow from the chromosphere into the corona. Variations in this region can influence the acceleration of the solar wind and play a role in the formation of prominences and other coronal structures. Observations of the transition region provide crucial constraints for theoretical models of atmospheric heating and solar wind generation.
Chromosphere
The chromosphere lies directly above the photosphere and extends several thousand kilometers upward. It is named for the reddish color it exhibits during solar eclipses, due to the strong H-alpha emission line of hydrogen. The temperature in the chromosphere ranges from about 4,500 K near the base to around 20,000 K at its upper boundary. This temperature gradient is surprising, as it increases with altitude, suggesting active heating processes not present in the photosphere.
The chromosphere is structured by magnetic fields that form spicules, fibrils, and plages—jet-like and bright structures observable in spectral lines such as H-alpha and Ca II K. These features are constantly changing and reveal the magnetic complexity of the solar atmosphere. Spicules, for example, are narrow columns of plasma that rise and fall on timescales of minutes, contributing to the mass and energy flow between the photosphere and corona. The chromosphere is the site of many transient events such as surges, Ellerman bombs, and microflares, all of which play roles in coronal heating theories.
Because the chromosphere is partially ionized, its plasma dynamics are governed by both gas pressure and magnetic forces. Wave heating, particularly from Alfv’en waves and acoustic shocks, is thought to contribute to the temperature rise observed in this layer. High-resolution observations from instruments like the Interface Region Imaging Spectrograph (IRIS) have revealed fine-scale structures and rapid dynamics that challenge existing models. The chromosphere is thus a highly complex and active region crucial to understanding stellar atmospheres.
Photosphere
The photosphere is the visible surface of the star and the layer from which most of the star’s light is emitted. It marks the boundary between the opaque interior and the transparent outer atmosphere. With temperatures averaging around 5,778 K, the photosphere radiates as a near-perfect blackbody, peaking in the visible portion of the electromagnetic spectrum. Sunlight observed from Earth originates in this layer.
The photosphere is not uniform but exhibits a granular texture caused by convection beneath the surface. Granules are small, bright cells of rising hot plasma surrounded by darker, cooler sinking material. Each granule spans about 1,000 km and lasts only a few minutes. These features are the surface manifestation of convective motions that transport energy from the interior to the surface. On larger scales, supergranulation patterns represent broader convective flows and magnetic field organization.
Magnetic fields in the photosphere are responsible for the formation of sunspots, plages, and faculae. These magnetic features affect radiative transfer and can locally suppress convection, causing temperature variations. Spectral lines formed in the photosphere provide critical diagnostics of temperature, pressure, and composition. Limb darkening, a phenomenon where the solar disk appears darker near the edges, is a direct consequence of the photosphere’s temperature gradient with depth.
Convective Zone
The convective zone is the outermost layer of the stellar interior, extending from the photosphere down to about 0.7 solar radii in the Sun. In this region, energy is transported primarily through convection rather than radiation. Hot plasma rises toward the surface, cools, and then sinks back down in a cyclic pattern. This convective motion is essential for the star’s thermal balance and influences the structure of the photosphere and chromosphere above.
Convection in this zone is driven by the increasing opacity of the plasma as it cools and becomes partially ionized. Radiative transfer becomes inefficient, and convection becomes the dominant transport mechanism. The convective zone is turbulent and dynamic, generating the surface granulation pattern and driving oscillations observable as helioseismic waves. These oscillations offer a powerful tool for probing the internal structure of the star.
The base of the convective zone is a region of intense scientific interest known as the tachocline, a shear layer where differential rotation between the radiative and convective zones generates magnetic fields through the solar dynamo process. These magnetic fields are thought to be responsible for the Sun’s magnetic cycle, including sunspot formation and reversal of the global magnetic field approximately every 11 years. The convective zone thus plays a key role in both energy transport and magnetic activity.
Radiative Zone
Beneath the convective zone lies the radiative zone, where energy produced in the core is transported outward by the process of radiative diffusion. This layer extends from roughly 0.25 to 0.7 solar radii and is characterized by a stable stratification, meaning that convection is suppressed due to the relatively efficient energy transport by photons. The temperature in this zone decreases gradually from over 7 million Kelvin near the core to about 2 million Kelvin at the tachocline.
In the radiative zone, photons are absorbed and re-emitted countless times in random directions, a process that results in an extremely long transit time for energy to reach the surface—on the order of thousands to millions of years. The high density and ionization levels ensure that radiation is the most efficient means of energy transport. The radiative zone is also relatively quiet and stable, making it an excellent environment for the propagation of pressure waves used in helioseismology.
Because the radiative zone is non-convective, it does not contribute directly to the magnetic dynamo. However, its interaction with the overlying convective zone at the tachocline is critical for magnetic field generation. The temperature gradient here is sub-adiabatic, meaning that displaced material tends to return to its original position rather than continue moving. This stability contributes to the long-term structural equilibrium of the star.
Core
The core is the central region of a main-sequence star where nuclear fusion occurs. In the Sun, it spans roughly the innermost 0.25 solar radii and contains about 50
The energy produced in the core is released as high-energy gamma photons, which gradually lose energy and scatter as they travel outward through the radiative and convective zones. This energy supports the star against gravitational collapse and maintains the hydrostatic equilibrium. The core’s temperature and density are precisely balanced to sustain fusion at a steady rate, regulating the star’s luminosity and lifetime.
The core evolves over time as hydrogen is converted into helium. As the hydrogen supply diminishes, the core contracts and heats up, eventually igniting new fusion processes if the star has sufficient mass. In lower-mass stars, this leads to red giant formation; in higher-mass stars, it initiates subsequent fusion cycles including helium burning and beyond. Thus, the core is not only the engine of stellar energy but also the cradle of nucleosynthesis and stellar evolution.
Flares
Solar flares are sudden, intense bursts of radiation emanating from the solar atmosphere, primarily near active regions associated with sunspots. These events occur when magnetic energy that has built up in the solar corona is suddenly released, often due to magnetic reconnection. Solar flares emit energy across the entire electromagnetic spectrum, from radio waves to gamma rays, and can last from minutes to hours depending on their magnitude.
The energy released during a flare can reach \(10^{25}\) Joules or more, accelerating charged particles and heating surrounding plasma to tens of millions of Kelvin. Flares are typically classified by their X-ray brightness in the 1-8 Angstrom range into categories such as A, B, C, M, and X. An X-class flare represents the most powerful type and can significantly affect Earth’s ionosphere, disrupting radio communication and GPS systems.
Flares are often associated with coronal mass ejections (CMEs), although the two phenomena are distinct. While flares release intense radiation, CMEs involve large-scale ejections of plasma. The precise relationship between flares and CMEs remains an active area of research. Flares provide critical insight into the workings of stellar magnetic fields and plasma physics, and their study has applications in both astrophysics and space weather forecasting.
Sunspots
Sunspots are cooler, darker regions on the photosphere caused by concentrated magnetic fields that inhibit convection. The strong magnetic field—often over 3,000 Gauss—reduces the efficiency of heat transport from the interior, leading to temperatures that are about 1,500 K cooler than the surrounding photosphere. Despite appearing dark, sunspots are still extremely bright and would shine intensely if isolated from the Sun.
Sunspots typically appear in pairs or clusters aligned along magnetic field lines. Each spot consists of a darker central region called the umbra, where the magnetic field is strongest and most vertical, and a lighter surrounding area called the penumbra, where the field is weaker and more inclined. The size of sunspots can range from a few thousand to tens of thousands of kilometers in diameter, and their lifetimes vary from days to several weeks.
Sunspot numbers follow an approximately 11-year cycle known as the solar cycle, during which the Sun’s overall magnetic field reverses polarity. The number and distribution of sunspots correlate with solar activity, including flares, prominences, and coronal mass ejections. Observing sunspots provides critical data for understanding the solar dynamo and predicting space weather. Their behavior also serves as a proxy for stellar magnetic activity in other Sun-like stars.
Stellar Classification
A reoccurring topic in our discussion of the properties of stars has been the idea of temperature and spectral classification. Given the sheer diversity of stellar composition and physicality, astronomers created a few key systems of classification to group certain stars together for easier identification and better understanding of their evolutionary relationships.
Harvard Classification System
The first and most widespread stellar classification system is the Harvard Classification system, which divides stars into letter categories from O to M, O being the hottest stars (and thus also most massive stars) while M stars are typically low massed and always low temperatures. With their differing masses and temperatures comes unique spectral and compostitional properties that we shall enumerate below. The following list (the classification system goes OBAFGKM in order of hottests to coldest) is from hottest to coldest type.
O Type Stars: O type stars are the most massive and hottest type of stars, typically at 30,000 K surface temperatures. Because they have such high temperatures, they are blue in color. Their sizes are typically greater than 6.6 R\(_\odot\) radii and their main sequence luminosities are greater than 30,000 L\(_\odot\). O type stars are masses typically greater than 16 M\(_\odot\). Regarding their spectral lines, O type stars show weak hydrogen lines, very strong ionized helium (He II), prominent ionized silicon (Si IV), ionized oxygen (O II), ionized nitrogen (N III), ionized carbon (C III), neutral helium lines (He I), and Balmer lines (in earlier types). Because of their high UV radiation, O type stars peak in the UV blackbody spectrum. Their B-V index is -0.33 (remember that the more negative, the hotter and the bluer the star).O type stars are relatively rare, encompassing 0.000013 percent of all stars in our universe
B Type Stars: B type stars typically around 20,000 K in surface temperature are a deep bluish white color. They range in size from 1.88 R\(_\odot\) to 6.6 R\(_\odot\). B type stars have main sequences luminosities from 25 to 30,000 L\(_\odot\) (yes that big of a range), and masses between 2.1 M\(_\odot\) to 16 M\(_\odot\). Regarding their emission lines, B type stars have medium/moderate Hydrogen lines in their spectra, strong neutral helium (He I), Balmer lines, neutral and ionized silicon (Si I and II), and neutral Oxygen (O I). These emission lines can be generalized to be high in "singly ionized heavy elements." Because of their high UV emission, the light of B type stars peak in the UV blackbody spectrum. B type stars have a B-V index of around -0.3 and account for approximately 0.13 percent of all stars in our universe.
A Type Stars: A type stars are typically around 10,000 K in surface temperature and white in color. They range in size from 1.4 R\(_\odot\) to 1.8 R\(_\odot\). A type stars have main sequences luminosities from 5 to 25 L\(_\odot\), and masses between 1.4 M\(_\odot\) to 2.1 M\(_\odot\). Regarding their emission lines, A type stars are the strongest in Hydrogen (neutral H) and Balmer Absorption lines out of any other star type, weak in ionized metals like ionized iron (Fe II), ionized magnesium (Mg II), and ionized calcium (Ca II, emitted at 3934 nm wavelengths), and have an absence of neutral Helium (unlike B and O stars). Like K type stars, A type stars also contain Si I in its emission spectrum. A type stars account for around 0.6 percent of all stars in the universe and have a B-V index of around -0.02
F Type Stars: F type stars are typically around 7000 K in surface temperature and yellowish-white in color. They range in size from 1.1 R\(_\odot\) to 1.4 R\(_\odot\). F type stars have main sequences luminosities from 1.5 to 5 L\(_\odot\), and masses between 1.04 M\(_\odot\) to 1.4 M\(_\odot\). Regarding their emission lines, F type stars have medium/moderate Hydrogen lines, strong in ionized calcium (Ca II), neutral iron (Fe I), neutral calcium, neutral cromium, and a few molecules like CH seen at 4300 nm wavelengths. F type stars account for around 3 percent of all stars in the universe and have a B-V index of around 0.3, signifying their cooler temperatures.
G Type Stars: G type stars are typically around 6000 K in surface temperature and yellow in color. Our sun is a G type star. They range in size from 0.9 R\(_\odot\) to 1.1 R\(_\odot\). F type stars have main sequences luminosities from 0.6 to 1.5 L\(_\odot\), and masses between 0.8 M\(_\odot\) to 1.04 M\(_\odot\). Regarding their emission lines, G type stars have weak Hydrogen lines, very strong in calcium (Ca I and II), strong sodium (Na I and II), other a neutral and ionized metals. G type stars account for around 7.6 percent of all stars in the universe and have a B-V index of around 0.58.
K Type Stars: K type stars are typically around 4000 K in surface temperature and light orange in color. They range in size from 0.7 R\(_\odot\) to 0.9 R\(_\odot\). F type stars have main sequences luminosities from 0.08 to 0.6 L\(_\odot\), and masses between 0.45 M\(_\odot\) to 0.8 M\(_\odot\). Regarding their emission lines, K type stars have weak Hydrogen lines, very strong in calcium (Ca I and II), strong sodium (Na I and II), other a neutral and ionized metals (yes just like G type stars). K type stars, unlike G type stars, are the strongest in neutral metals and most likely to be habitable. Along with M type stars, K type stars can take the form of a red/sub giant star. the K type stars account for around 12.1 percent of all stars in the universe and have a B-V index of around 0.81.
M Type Stars: M type stars are the coolest stellar type, typically ranging around 3000 K in surface temperature and orangish red in color. M stars have sizes less than 0.7 R\(_\odot\). M type stars have main sequences luminosities less than or equal to 0.08 L\(_\odot\), and masses between 0.08 M\(_\odot\) to 0.45 M\(_\odot\). Regarding their emission lines, M type stars have very weak Hydrogen lines, strong in molecular titanium oxide (titanium oxide is a distinctive feature of M stars that you should memorize), cyanogen (CN), and other oxides. As you could’ve reasonable deduced, M stars are unique in their compsition of molecules in their spectra. This is especially true for the hydrogen in its spectra, which is now molecular form (H2). Hydrogen is thus weak in their emission lines because H2 molecules are non-polar and hard to detect, though they are high concentration in M stars. The molecular composition of M stars is a direct result of their low temperatures. M type stars account for around 76.5 percent of all stars in the universe and have a B-V index of around 0.58. M type stars are the most common stars in the universe. The specific type of M star that is the most common star is called a red dwarf.
In addition to these 7 classes there are three more that are below M in temperature and are used to classify different types of brown dwarf stars. These are L, T, and Y type stars. Here is what is unique about each additional classification:
L Type Stars: L stars range from 1300 to 2400 K in surface temperature and are the color red. They have an atmosphere dominated by carbon monoxide, and have emission lines strong in Metal hydrides and alkali metals (Na, K, Rb, etc).
T Type Stars: T stars range from 700 to 1300 K in surface temperature and are a deep magenta in color. T stars have a Methane dominated atmosphere and strong methane emission lines.
Y Type Stars: Y stars are the coolest type of star on the Harvard classification system, typically maintaining less than 700 K in surface temperature. They are deep red (almost a brown) in color, and are very similar to cloudy gas giant planets. Y type stars have distinctive ammonia lines in their spectra.
Yerkes Classification & Arabic Numerical System
Two common classification systems used together with the Harvard classification system are the Yerkes Classification system and the Arabic Numerical system.
The Yerkes classification system adds a another layer of classification to the Harvard classification system, allowing astronomers to group stars by evolutionary track/stage as well as temperature. Here are the subclassifications of the Yerkes system and what they mean:
0 or Ia^+: hypergiants or extremely luminous supergiants.
Ia: luminous supergiants
Iab: intermediate-size luminous supergiants
Ib: less luminous supergiants
II: bright giants
III: normal (red) giants
IV: subgiants
V: main sequence stars
sd or VI: subdwarfs
D or VII: white dwarfs, neutron stars
Stars in the process of transitioning between two different evolutionary stages can be designated with a dash in between the two roman numerals (ex. BIV-V would be a B star transitioning from the main sequence stage to being a subgiant).
The Arabic Numerical system is significantly simpler than both the Harvard and Yerkes classification systems as it simply is a number from 0 to 9 designating the temperature magnitude of a certain star, given the temperature of the stars spectral type (OBAFGKM, that is). 9 represents the coolest of the stars of a certain type and 0 is the hottest. An example of this is a O9II star, which is the coolest type of O star which is a bright giant. Easy, right?