![\"how](\"https:\/\/www.newsbeep.com\/nz\/wp-content\/uploads\/2026\/02\/https___specials-images.forbesimg.com_imageserve_5f10f05db8201200063cd6b1_The-open-star-cluster-NGC-.jpeg\")
<\/p>\nIf there\u2019s one thing we can be certain of when we look out at the glittering canopy of the night sky, it\u2019s this: that someday, all of those luminous points of light, including every star and every galaxy, will someday fade away and cease to shine. The stars and stellar remnants, the primary sources of light and heat and energy that propagate throughout the Universe, are powered by finite sources of fuel: whether through nuclear fusion, gravitation, or any other mechanism. At some point, those fuel sources will be exhausted, no further energy will be naturally extracted from what remains within them, and those once-brilliant objects will fade away into darkness. Some stars live only briefly, others will continue to shine long into the future, with lifetimes far exceeding our Universe\u2019s current 13.8 billion year age.<\/p>\n
That brings us to the question of James D, who was curious about the longest-lived stars of all, and wrote in to ask:<\/p>\n
\u201cI was reading one of your articles about the lifespan of red dwarf stars, with the smallest living anywhere from 20 trillion to 380 trillion years as a theoretical limit. There isn\u2019t very much information I can find on the internet about what influences the lifespan, so I was wondering how much metallicity affects the lifespan of the star, and what factors overall play into the theoretical limit for red dwarf longevity?\u201d<\/p>\n
James, in an era dominated by AI slop, you did the right thing by coming to a human with actual, expert-level knowledge. Let\u2019s walk you \u2014 and everyone else \u2014 through the science behind the longest-lived stars of all.<\/p>\n
<\/p>\n
This Hubble Space Telescope image of open star cluster NGC 290, showcases a region where thousands of newborn stars were created 30-60 million years ago. They come in a wide variety of masses, where a combination of their initial mass and future interactions will determine their ultimate fates.\n<\/p>\n
Credit<\/a>: ESA and NASA; Acknowledgment: E. Olszewski (University of Arizona)<\/p>\n When a new star is born, if it remains in isolation (i.e., doesn\u2019t merge or interact with any other massive objects, like other stars), nearly everything about its future history can be calculated. Like anything physical, a star\u2019s physical properties are primarily determined by its composition. Its mass is the most important factor in determining both its lifetime and its fate, with other secondary factors, such as metallicity (or the fraction of heavy elements present within it), also playing a role. Additionally, in order to shine, stars need to be hot: hot enough to ignite nuclear fusion in their cores. Without that key step, you can\u2019t be classified as a true star; the presence of those fusion reactions, where hydrogen gets fused into helium, separates stars from all other heavenly bodies. <\/p>\n Early on in cosmic history, the Universe was composed primarily of hydrogen and helium, as no substantial quantities of heavier elements were formed during the early stages of the hot Big Bang. Once stars begin to form, the Universe becomes \u2014 depending on your perspective \u2014 either:<\/p>\n enriched, as the heavy elements formed via nuclear fusion (and other nuclear processes) within stars and stellar cataclysms get returned to the interstellar medium and participate in future generations of star-formation,<\/p>\n or polluted, as those same heavy elements \u201ctaint\u201d the pristine hydrogen and helium reservoirs, and everything that forms out of them, in subsequent generations of processes.<\/p>\n Either way, that\u2019s the material we begin with, at any epoch throughout cosmic history, when the Universe forms stars.<\/p>\n This nebula in the Perseus molecular cloud, NGC 1333, is located only 960 light-years away here in our own Milky Way. While Hubble can only capture the light-blocking dust and heated gaseous material, JWST is spectacular at viewing an enormous number of obscured stars and protostars inside as well as the cooler material that is heated by the environmental conditions, which is invisible to Hubble.\n<\/p>\n Credit<\/a>: NASA, ESA, STScI<\/p>\n From a sufficiently massive, sufficiently cold cloud of material, whether pristine or enriched, gravitational collapse will begin to occur. (If the particles within the cloud are moving too fast for the amount of mass that\u2019s present, collapse won\u2019t occur; it will remain a significantly large-sized cloud.) As the overdense regions within the cloud attract more and more matter, they won\u2019t just grow into massive clumps, those clumps will also trap the heat generated from gravitational collapse, causing the interior of the clump to heat up. That leads to high temperatures, which create a glowing protostar due to the rapid (kinetic) motion of the internal gas particles, and eventually, after a few tens of millions of years, the core temperatures rise high enough (above 4 million K or so) that nuclear fusion of hydrogen begins.<\/p>\n It\u2019s that moment \u2014 the moment that the first nuclear reactions begin that fuse hydrogen nuclei (protons) into helium nuclei \u2014 that signals the true \u201cbirth\u201d of a star. For the hottest, most massive stars, the primary pathway by which hydrogen fusion occurs is through what\u2019s called the C-N-O cycle<\/a>: where hydrogen nuclei get added into pre-existing heavier (carbon, nitrogen, and oxygen) nuclei, liberating energy and leading to the eventual production of helium atoms, restoring a new carbon nucleus at the end of each iteration of the cycle. Meanwhile, for lower mass stars, including stars like our Sun, it\u2019s the proton-proton chain<\/a> that dominates, where hydrogen fuses into helium, and liberates energy, primarily through that mechanism instead.<\/p>\n This illustration of the lowest-energy component of the CNO cycle, which is the most common mechanism by which it occurs in the Sun, details how hydrogen fuses into helium as a result of chain reactions involving carbon, nitrogen, and oxygen. In stars with more than 130% the mass of the Sun, this, rather than the proton-proton chain, dominates as far as nuclear fusion is concerned.\n<\/p>\n Credit<\/a>: Borb\/Wikimedia Commons<\/p>\n This makes sense from a certain perspective. If you want to bring two tiny, like-charged particles together, you absolutely have to have enough energy so that those particles can approach one another closely enough that they can \u2014 for lack of a better term \u2014 \u201ctouch\u201d one another. Because these atomic nuclei are all positively charged, and also because they are quantum systems, we can consider them as waves (with wavefunctions) instead of simply as classical point particles, what defines whether they \u201ctouch\u201d or not is whether their quantum wavefunctions overlap or not. If they do overlap, then there\u2019s the potential for a fusion reaction to occur. If not, then there\u2019s no chance at all.<\/p>\n Inside a star\u2019s core, or the region where fusion reactions occur, the opportunities for particle overlap depend on how fast two colliding particles are moving when they smash into one another. Although that depends on temperature severely, it\u2019s also an inherently quantum process: one that requires particles to tunnel into a more stable state during that brief period when their wavefunctions overlap. Hotter stars, therefore, have:<\/p>\n higher rates of fusion,<\/p>\n larger regions of their core where fusion occurs,<\/p>\n and more opportunities for fusion to proceed through greater numbers of pathways,<\/p>\n whereas stars with cooler cores have small regions where fusion occurs, slow rates of fusion, and can only experience the proton-proton chain. The more massive your star, the higher its core temperature; the less massive your star, the lower its core temperature.<\/p>\n When two protons meet each other in the Sun, their wavefunctions overlap, allowing the temporary creation of helium-2: a diproton. Almost always, it simply splits back into two protons, but on very rare occasions, a stable deuteron (hydrogen-2) is produced, due to both quantum tunneling and the weak interaction.\n<\/p>\n Credit<\/a>: E. Siegel<\/p>\n This leads to a fascinating but sometimes counterintuitive conclusion: that the longest-lived stars are actually going to be the ones with the least total amount of fuel in them. Remember, the most massive stars, or the ones with the most amount of fuel in them, are going to have the largest, hottest cores, and within those cores, they\u2019ll have the fastest rates of fusion, the greatest number of realized fusion pathways inside of them, and the greatest energy outputs. The most massive stars are going to be hotter, larger, bluer, more luminous, but also shorter-lived than the less massive stars.<\/p>\n That\u2019s heavily related to what their core temperatures are. Whereas fusion (of hydrogen into helium) begins when core temperatures reach 4 million K, our Sun\u2019s core gets up to about 15 million K. At those temperatures, a yellow star like the Sun gets about 1% of its energy from the CNO cycle, whereas a low-mass red dwarf only gets energy from the proton-proton chain. The rate of fusion also affects a star\u2019s equilibrium temperature, which impacts what we see as far as \u201ccolor temperature,\u201d or the temperature of the star at the edge of its photosphere, looks like. The lowest-mass stars appear small, faint, red, and cool; the higher mass stars appear large, bright, blue, and hot.<\/p>\n The (modern) Morgan\u2013Keenan spectral classification system, with the temperature range of each star class shown above it, in kelvin. In terms of size, the smallest M-class stars are still about 12% the diameter of the Sun, but the largest main sequence stars can be dozens of times the Sun\u2019s size, with evolved red supergiants (not shown) reaching hundreds or even 1000+ times the size of the Sun. A star\u2019s (main sequence) lifetime, color, temperature, and luminosity are all primarily determined by a single property: mass, with other properties (like metallicity) playing only minor, secondary roles.\n<\/p>\n Credit<\/a>: LucasVB\/Wikimedia Commons; Annotations: E. Siegel<\/p>\n However, one of the more annoying facts about our Universe is that we can only see it as it is now: 13.8 billion years after the Big Bang. Sure, we can look to great distances and see the Universe as it was when it was younger; given that the Universe is huge and light only travels at a finite speed, we can see objects as they were when their arriving-now light was first emitted, even if that\u2019s from millions or billions of years ago. But if there are objects that form whose lifetimes are longer than the lifetime of the Universe at present, we won\u2019t have even a hope of observing how these stars move through the full stages of their stellar life cycles.<\/p>\n That\u2019s why, when we talk about how the mass of a star relates to the lifetime of a star, we can\u2019t exactly measure that directly for stars whose lifetime is longer than the present age of the Universe. A star that\u2019s maybe 70-80% of the Sun\u2019s mass would \u2014 if it was born at the first moment of the hot Big Bang \u2014 just be reaching the end of its life now. That\u2019s the limit on how massive a star can be where we can still observe it in different stages of stellar evolution. But less massive stars, and stars can have as little at 7.5% of the mass of the Sun and still be stars, are still in that first, long stage of their life: where they spend time on the main sequence, burning through their nuclear fuel in their cores and fusing hydrogen into helium, primarily through the proton-proton chain.<\/p>\n This color-magnitude (or Hertzsprung-Russell) diagram shows a \u201csnapshot\u201d of color vs. magnitude of a wide variety of stars. When stars ignite nuclear fusion in their cores for the first time, they begin life at the bottom of the main sequence (vertically) for whatever their color is. Over their hydrogen-burning lifetimes, they migrate upward, becoming brighter but remaining at approximately the same color\/temperature, before they run out of hydrogen in their cores and begin evolving first into subgiants, and then into red giants or supergiants, where they then head into the final stages of their lives and approach their stellar demises.\n<\/p>\n Credit<\/a>: Richard Powell\/Wikimedia Commons<\/p>\n Therefore, it\u2019s only for the higher-mass stars that we can see their evolution:<\/p>\n into red giants,<\/p>\n into the asymptotic giant branch phase,<\/p>\n and into the final stages of their lives (dying in supernovae or planetary nebulae),<\/p>\n and potentially leaving remnants (like white dwarfs, neutron stars, or black holes) behind.<\/p>\n By observing those latter stages of evolution, we can learn quite a bit about the impact of secondary effects, like their metallicities, on their overall lifespans. This has been an active area of study for some time, and what we\u2019ve learned is the following.<\/p>\n Initially, the presence of greater amounts of heavy elements behaves as though there\u2019s an extra absorptive presence in the star\u2019s interior, acting to increase the star\u2019s opacity, making it less transparent to photons, and also inhibiting potential fusion reactions<\/a>. This slows down the star\u2019s rate of fusion, meaning that it takes longer to burn through the same amount of fuel. However, this is most easily observed in the most luminous and\/or the most massive stars, such as stars in the asymptotic giant branch phase. Those stars with higher metal contents live longer and burn cooler in general, but only slightly, and in a way that\u2019s much more impactful for higher mass stars than lower-mass stars.<\/p>\n This three-track graph shows the mass of stars found in post-main-sequence phases as a function of mass and metallicity\/ For stars of the same mass, higher-metallicity stars live for longer, while lower-metallicity stars live for shorter amounts of time. However, this analysis is only good for stars down to about 0.8 solar masses; below that, the Universe isn\u2019t old enough for us to have any meaningful data.\n<\/p>\n Credit<\/a>: E.M. Manning & A.A. Cole, Monthly Notices of the Royal Astronomical Society, 2017<\/p>\n For these shorter-lived classes of stars, including the Sun-like stars and all more massive stars, we can actually observe their lifetimes end-to-end, as well as model what\u2019s going on in their interiors. One important lesson that we learn is that, for these stars, the material that fuses inside of them \u2014 i.e., the material that\u2019s in their cores \u2014 remains in their cores during the entire life cycle of the star. There is no efficient transport of material in-and-out of the core, and thus, most of the star\u2019s interior never gets a chance to fuse simply because of its location. What doesn\u2019t start in the core never encounters those high-enough temperatures for fusion to occur, what starts in the core doesn\u2019t leave the core, ensuring that the \u201cnext stage\u201d in stellar evolution arrives before fusion ever runs to full completion.<\/p>\n But for the lowest-mass stars, they have the smallest cores, the lowest temperatures of any true stars (from the core to the photosphere), the slowest rates of fusion, and the longest timescales for burning. This means that processes that simply \u201ctake too long\u201d for us to observe in our 13.8 billion year old Universe can occur, and eventually will. That includes the process of whole-star convection, where fused material in the core gets transported out of the core, and where unfused, raw material outside of the core gets transported into the core. This low-and-slow process of nuclear fusion in the lowest mass stars ensures that 100% of the raw hydrogen inside of them will eventually burn, fusing to completion.<\/p>\n This photo showcases Proxima Centauri: the closest star to our own Sun at present. Although it\u2019s only 4.24 light-years away, Proxima Centauri is not even close to visible to the naked eye, as it\u2019s intrinsically nearly 1000 times fainter than the Sun: typical of red dwarfs, the most common but faintest type of star in the Universe.\n<\/p>\n Credit<\/a>: Alessandro Cipolat Bares<\/p>\n This would seem to imply that a \u201cpolluted\u201d star at the low-mass end of what becomes a red dwarf has a disadvantage compared to a \u201cpristine\u201d star at the low-mass end of red dwarfs. Because there\u2019s a fraction of that star that\u2019s already been burned and fused, that means there\u2019s a deficiency of material that can be further fused inside of it, and hence you\u2019d think that means it would have a shorter lifetime.<\/p>\n But there\u2019s a counterargument to that. Perhaps the increased metallicity means that the opacity inside the polluted star\u2019s core will be greater compared to the pristine star, giving it an even lower rate of fusion and an even smaller energy output per-unit-time, and hence, a longer lifetime. Perhaps, on extremely long timescales, those heavier elements will sink to the center of the core, producing an inert center that slows the rate of fusion even further.<\/p>\n Depending on which effect wins \u2014 and, unfortunately, we have no data on this, only our theoretical calculations and the sandbox of our simulations \u2014 the lifetime of a higher metallicity low-mass red dwarf could be slightly greater or slightly less than the lifetime of a more pristine low-mass red dwarf of equivalent initial mass.<\/p>\n This graph shows the brightness of a star (y-axis) versus the mass of a star (x-axis) over the course of its main sequence lifetime. Less massive stars get much fainter (as a function of mass) very swiftly, proportional to the mass ratio to the 5th power, while the most massive stars increase in luminosity more slowly, as mass to the 4\/3 power. The gradual change in the relationship is shown by the blue curve.\n<\/p>\n Credit: Kirk Korista; Private communication<\/p>\n However, it\u2019s worth pointing out that there are enough uncertainties just with calculating \u201cthe lifetime of a low-mass red dwarf star based solely on its mass\u201d that the metallicity corrections are small compared to them. Consider Proxima Centauri<\/a>, for instance: the closest star to the Sun at present, at 4.24 light-years away, but a star that\u2019s only 12% the mass of our Sun. It is so faint that:<\/p>\n visually, it would take about 20,000 of them together to be as intrinsically bright as our Sun,<\/p>\n energetically, it would take about 640 Proxima Centauris to equal the energy output of one Sun,<\/p>\n but as far as its lifetime goes, Proxima Centauri will no doubt live for many hundreds, and perhaps thousands, of times as long as our Sun will.<\/p>\n We don\u2019t know how the lifetime of a red dwarf changes, or how or if the rate of fusion slows, when it nears the end of its life and has burned through more of its fuel. We don\u2019t know whether having heavier elements (carbon, oxygen, iron, etc.) in a red dwarf\u2019s core alters the rate of fusion compared with having large amounts of helium: the end product of the proton-proton chain. We don\u2019t know how slow the rate of fusion is in the absolute lowest-mass red dwarf star possible, and whether that rate is stable over time or not. All we can say, for certain, is that there will be red dwarf stars that continue to shine for tens of trillions of years and probably for over 100 trillion years. My original 20-to-380 trillion estimate<\/a> is just that: an estimate, and while you might wiggle that number by a few trillion years in either direction due to our unknowns about metallicity, the uncertainties about red dwarf lifetimes solely based on mass are still far greater than that!<\/p>\n![\"A](\"https:\/\/commons.wikimedia.org\/wiki\/File:NGC290.jpg\")
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