- Astronomical Dossier: The Sun (at a Glance)
- Nuclear Fusion in the Sun’s Core
- Lifespan of a Solar Gamma Ray (1,000,000 Years + 8 Minutes and 20 Seconds)
- Atmosphere of the Sun: The Photosphere, Chromosphere & Corona
- The Sun’s Magnetic Fields and Sunspots
- Prominences, Flares, CMEs and the Solar Wind
- Bibliography & References
- Photo Credits
Sol: Giver of Life
“Sol,” which is the Latin word for Sun, isn’t the biggest or most impressive star around. Although we cannot perceive our stalwart patron at its apex without damaging our eyes – due to its brilliantly blinding apparent visual magnitude (mv) of -26.74 – were the Earth located just ten parsecs away the star would be barely perceivable to us in the night sky, with an absolute visual magnitude (Mv) of 4.83. Even so, Sol is the sole supplier of energy for our planet.
For example, think back to your last meal and consider this: no matter what you had to eat and/or drink, you were dining on sunlight. Moreover, you weren’t the only one – every living thing on the planet is dependent on our local star for survival; specifically, from the steady stream of energy it provides us in the form of electromagnetic radiation (e.g., light). Plants and plankton convert sunlight energy directly through the process of photosynthesis, and animals consume that energy indirectly by eating plants and plankton (or by eating other animals that eat plants and plankton).
The Earth itself was born of, molded, and brought to life by the Sun. According to current theory, gas, dust and other debris left over from the interstellar medium nebula where the Sun was born condensed into orbiting, colliding bodies that (eventually) arranged themselves into a handful of protoplanets. Around four-and-a-half billion years ago, one of the more rocky of these lifeless chunks had managed to assimilate most of the larger debris in its orbit, and was busily arranging the various metals comprised therein into (what would later become) our planet’s core, mantle and crust.
(Note: In Astronomy, every element besides hydrogen or helium is referred to as a “metal.”)
Over the course of about a hundred million years, energy and gases released via Earth’s internal processes, working in tandem with the Sun’s unfailing illumination, allowed the oceans, the first continent-like landmasses and the atmosphere to form. (Readers may be surprised to learn that the formation of our atmosphere was thanks in large part to the Greenhouse Effect– but more on that in another lesson.) In any case, the winds followed suit as a direct consequence of temperature differentials in atmospheric gases caused by sunlight.
The Sun also plays an important ongoing role in determining the behavior of water here on Earth. From the tiniest of raindrops falling onto the gentlest of brooks resting lazily in an unnamed forest glen, to the raging rapids of the Corhu River in Turkey, or the vast and formidable network of Great Lakes in North America; from man-made marvels such as the Hoover Dam and the Panama Canal even unto all of the water in all of the oceans, the Sun is instrumental in driving our planet’s hydrologic cycle (i.e., the repeating process of evaporation, condensation and precipitation of Earth’s water vapor). In fact, it would not be an exaggeration to state that the cycle depends on sunlight! Moreover, even the tides are partly influenced by interactions with the Sun’s gravitational force.
To put it plainly: Every facet of our planetary biosphere evolved as a consequence of Earth’s link to the Sun.
Modern technology would not function without the star’s influence. The majority of our planes, trains and automobiles run on fossil fuels, which are merely the remains of plants and animals that lived hundreds of millions of years ago (thanks to the energy provided to them directly or indirectly in the form of sunlight). Similarly, the electricity that most of the world accesses by plugging computers, televisions, appliances, etc., into wall outlets is generated by burning fossil fuels. More environmentally conscious people might employ alternative energy technologies instead – e.g., solar, hydroelectric, wind – and these also depend on sunlight and/or the Earth’s interaction with the Sun.
In every way imaginable, we owe our lives to Sol. Whether you believe that our great benefactor was put in place by divine forces or not makes no difference to the humbling reality of our situation: we are children of the Sun; without it we would never have been here. Furthermore, we would not survive were the star suddenly to be removed from the sky. Just as energy is a prerequisite for reality as we know it, the Sun’s continuous supply of (electromagnetic radiation) energy is the prerequisite for our reality.
But, how does the Sun produce the energy that we need? For that matter, how long will it continue to produce that energy for us? What is the Sun made of? What are the physical, geographic and atmospheric characteristics of our trusty star, and how does it rate against other stars?
The answers to these questions, and more, will be explored in the following paragraphs.
Astronomical Dossier: The Sun (at a Glance)
The Sun is a luminous, fusion-powered, spherical celestial body made entirely of hot gas – that is, it’s a star. With a total mass of 1.989 x 1030 kg and a radius of 6.9599 x 105 km, it is average in size and luminosity (as compared to other main sequence stars), and it has a spectral classification of G2 V. Composed almost entirely of hydrogen and helium, with a small amount of trace elements, the Sun has six layers, each with distinct temperature, density and behavioral characteristics. From innermost to outermost, these are: the core, the radiative and convective zones, the photosphere, the chromosphere and the corona.
Our star has lived approximately five billion years out of its projected ten billion year life-cycle, and it may be accurately described as “middle-aged.” Earth maintains an average distance of 149,597,900 km (92,956,056.52 miles) from the Sun – a measurement that is also known as one astronomical unit (AU) – and while our planet makes one complete revolution every twenty-four hours, the Sun has a rotational period of 25.38 days (at the equatorial region of its photosphere). Although the densities of solar gases differ between layers, the Sun’s average density is 1.409g/cm3. Moreover, though the Sun’s temperature is not homogeneous between layers, its core and surface (i.e., photosphere) average temperatures are 1.5 x 107 kelvins (k) and 5.8 x 103 k, respectively.
Were a remarkably sturdy version of the NASA space shuttle to fly in too close to the Sun and get pulled down to its photosphere by gravity, the imperiled ship would need to achieve a velocity of 617.7 km/s before it could escape again. To place that in appreciable context, consider that the space shuttle need only achieve 11.2 km/s to escape Earth’s atmosphere! That might sound a little far-fetched until you are able to grasp just how large the Sun is from our frame of reference. But think of it this way: 99.8% of the total mass in our entire solar system is contained within the Sun, and you would need to string 109 Earths together in a line to equal the star’s diameter – i.e., the diameter of the Sun divided by the diameter of Earth equals the number of Earth’s needed, or:
(1,391,980 km)/(12756.2 km) = 109.1.*
Nuclear Fusion in the Sun’s Core
If we had technology that would allow us to withstand intense heat and pressure, we could take a journey 695,990 kilometers (approx. 432,561 miles) beneath the Sun’s photosphere, into its core. Here, we would find that weight from the star’s outer layers has compressed and accelerated the stellar core gas to blinding speeds, which also serves to raise the average core temperature to around 15,000,000 k. At this temperature the gaseous, completely ionized hydrogen particles (free protons) within the Sun’s core crash into each other with enough force to overcome the Coulomb barrier and trigger hydrogen fusion to form helium, in a step-by-step process known as the proton-proton chain (P-PC).
The reaction sequence comprising the P-PC is as follows:
1H + 1H = 2H + e+ + n
2H + 1H = 3He + g
3He + 3He = 4He + 1H + 1H
(Where “e+” is a positron, “n” is a neutrino and “g” is a gamma ray)
In the first reaction, ionized hydrogen particles (i.e., single protons) collide, yielding a deuterium nucleus, a positron and a neutrino. Next, the deuterium nucleus collides with another, single hydrogen particle to produce the 3He (i.e., light helium) isotope and a gamma ray. Finally, two light helium nuclei collide, yielding a common helium (i.e., 4He) nucleus and two protons (i.e., hydrogen nuclei). Note that in order for one common helium nucleus to be produced, the first two steps in the P-PC must occur twice.
Helium nuclei are created within the core at a rate of 1038 times per second, releasing a total energy of 1.365 x 1018 Joules/second. This may be easily calculated using Einstein’s formula, e = mc2, and simple arithmetic:
- The mass of four hydrogen nuclei/protons is: 4(1.67262178 × 10-27 kg) = 6.69048712 x 10-27 kg
- The mass of one common helium nucleus is equal to the mass of two protons and two neutrons minus the energy that must be released to bind these particles together:
2[(1.67262178 x 10-27 kg) + (1.674927352 x 10−27 kg)] – 28.3[(1.780 x 10-30 kg)] = 6.645 x 10-27 kg**
- By subtracting the mass of a common helium nucleus from the total mass of four hydrogen nuclei/protons, we see that the common helium nucleus is “lighter” than the sum of the masses of the four protons:
(6.69048712 x 10-27 kg) – (6.645 x 10-27 kg) = 4.549 x 10-29 kg
- Since energy can neither be created nor destroyed (only changed), we know that the “missing” 4.549 x 10-29 kg’s must have been converted into energy during fusion of the common helium nucleus. Using e = mc2, we may quantify that energy:
E = (4.549 x 10-29 kg)(3 x 108 m/s)2
E = 1.365 x 10-20 J
- Finally, we multiply the energy yield of one common helium nucleus fusion reaction (via the P-PC) by the number of these reactions that occur in the Sun per second, or:
(1.365 x 10-20J)(1038/s) = 1.365 x 1018 J/s
While approximately two percent of the energy released during fusion is carried off by neutrinos, the rest of it – in the form of gamma rays, positrons and the repellant motion of the “spare” protons (i.e., the protons that are expelled at the conclusion of the P-PC sequence) – serves to help keep the core gas hot. The gamma rays get absorbed almost immediately by the surrounding gas, and the positrons combine with free electrons to form gamma rays that are then absorbed by the surrounding gas, as well. When the errant protons left over from helium fusion reactions speed apart (i.e., due to their like, positive charges), they collide with other particles, and this also helps to heat the gas.
This process of sustained nuclear fusion within the Sun’s core, which has been ongoing for around five billion years, is the source of energy that allows all life to flourish on Earth. Specifically, the energy that we receive comes from gamma rays that are emitted as a consequence of the fusion process. However, the neatly packaged gamma ray photons of energy that are released in the core still have a long, transformative journey in front of them before they are ready to make the relatively short voyage to Earth.
Lifespan of a Solar Gamma Ray (1,000,000 Years + 8 Minutes and 20 Seconds)
As gamma rays that are floating around in the Sun’s core run into wayward electrons, they get deflected in random directions. This happens many times, and the gamma rays bounce around haphazardly on chaotic trajectories that invariably lead them out of the core. The gamma rays are then absorbed by gas in the layer of the Sun that is adjacent to the core – the cooler radiative zone, and are re-emitted as photons of smaller wavelengths (x-rays, ultraviolet rays, visible spectrum rays, etc.) over and over again, while simultaneously continuing to travel (or “radiate”) outward towards the photosphere. After tens of thousands of years, longer-wavelength photons that began as gamma rays in the core reach the Sun’s convective zone, which is an area directly beneath the photosphere where currents of hot gas rise and cooler gas sinks.
To understand why this happens, recall that the core’s average temperature is 15,000,000 k. But the temperature of the Sun’s interior recedes as distance from the core increases, and at 695,990 kilometers from the center (1 solar radius), the photosphere has a temperature of only 5,800 k. At these relatively low temperatures, the gas is less transparent to radiation than the deeper, hotter layers. As a result, energy gets backed up like a large cluster of marathon runners filtering through a narrow passage.
Convection occurs as hotter (more energetic) gas expands and rises, while cooler (less energetic) gas simultaneously contracts and sinks in a circulating pattern. That is, after warming convection gases into a rising, expanding state, energy is radiated out into space upon reaching the photosphere. This cools the gas, causing it to sink below the photosphere and absorb more energy in order to rise again on the next cycle, and so on.
Putting this all in perspective, by the time that a neatly bundled packet of energy that began life as a single gamma ray makes its trek through the core, the radiative and convective zones, then finally beyond the photosphere and out into space, it has been broken down into 1800 smaller photons. Moreover, the entire process can take as long as a million years. Contrarily, the voyage from the photosphere to Earth takes eight minutes and twenty seconds for the photons. This means that after having worked their way up from the intense, hellish conditions inside the Sun (over a period of time spanning more years than the entire history of the human race), the multiple fragments of a long-enduring gamma ray are finally able to experience the wonders of space travel for what must seem a heartbeat, before arriving at our little blue planet where they are subsequently absorbed by atmospheric molecules, plant cells or human retinas with little to no fanfare!
Atmosphere of the Sun: The Photosphere, Chromosphere & Corona
The Sun’s atmosphere consists of three layers: the photosphere, chromosphere and corona. Among these, only the photosphere is observable to the naked eye without the use of special equipment, except in the singular event of a total solar eclipse. In such instances, the chromosphere is briefly visible as a pinkish ring outlining the photosphere (just as the disk of the Moon completely covers the Sun), and the lower-corona appears as shining, nebulous white light surrounding the Sun like a halo.
Though the photosphere is only 500 km deep, most of the light that we receive from the Sun is radiated from this layer. Gas at the photospheric level maintains just the right average density (1 x 10-9g/cm3) to emit photonic energy, without becoming so dense that the energy gets trapped. Contrarily, light radiated from the inner layers of the Sun, though considerable, cannot penetrate the outer layers. Furthermore, the gases of the chromosphere and corona are not dense enough to emit a great deal of light, and what light they do emit is drowned out by the brilliant photosphere.
High-resolution images show the effects of convective zone activity on the photosphere in the form of granules, which are large, adjacent, cell-like regions with warmer (rising) gas in their centers and cooler (sinking) gas around their edges. Since granules are hotter in the center, their cooler edges appear dark in images – in accordance with the usual behavior of black body objects (i.e., the Stefan–Boltzmann law) – making the overall appearance of the photosphere grainy and mottled. With diameters of one or two thousand kilometers, granules completely cover the photosphere (known as granulation). Each cell lasts for ten to twenty minutes before dissipating, whereby new granules appear to replace the old. Specifically, the gases within granules have average rising/sinking velocities of 0.4 km/s.
Larger networks of granules have also been observed. Named supergranules, these phenomena contain around three-hundred granules each, and they have diameters approximately twice that of Earth. Supergranules appear to be caused by larger currents of slowly rising gas underneath the photosphere that last for one or two days.
The chromosphere is the second layer of the Sun’s atmosphere. It has an irregular form, with peaks and valleys of ionized gas that resemble jagged, undulating spires on an otherwise flat plain (i.e., the “plain” of the photosphere). Curiously, the temperature of these spires (known as spicules) – and of the rest of the lower-chromospheric gas – is cooler than you might expect. In fact, the chromosphere’s average temperature drops to around 4500 k just above the photosphere, before rising drastically as altitude increases. At the upper-most level of the chromosphere, just below the lower corona, the gas reaches temperatures of 1 x 106k.
Whereas the photosphere produces an absorption spectrum, the chromosphere produces an emission spectrum. According to Kirchhoff’s second law, this informs us that the gases of the chromosphere must have a low-density – in this case, 1 x 10-12 g/cm3 (or, about three thousand times less dense than Earth’s air at sea level). Based on spectroscopic analysis, astronomers have been able to determine that the chromosphere’s pinkish color, as seen during a solar eclipse, is produced by the merging of three visual wavelength Balmer lines that are characteristic to excited hydrogen: red, blue and violet.
However, at higher altitudes where gas is even more ionized, astronomers have discovered that the chromosphere produces dark lines on photospheric absorption spectra at certain wavelengths.*** When viewed at such wavelengths (in images known as filltergrams), the chromosphere’s gaseous spicules may be seen rising above the edges of supergranules like wisps of grass growing around a poorly-maintained sidewalk. By analyzing filtergrams and short-wavelength solar images (e.g., ultraviolet or x-ray), astronomers are better able to study the higher regions of the chromosphere.
The third and final layer of the Sun’s atmosphere is also, perhaps, the most enigmatic. Like the chromosphere, it is completely invisible to human eyes, except during a solar eclipse when its lower region may be viewed as a white, mist-like substance surrounding the disk of the Sun. But using special telescopes called coronagraphs, astronomers have been able to measure the corona out beyond 13,919,800 km (around 8,649,360 miles).
The gases in the corona have an extremely low density: 106 atoms/cm3 in the lower levels, and 1 to 10 atoms/cm3 in the higher levels; therefore, not much light is radiated from this region. Some light emitted from the corona has the same spectral properties as the absorption spectra of the photosphere, but this is just photospheric light reflected off of dust particles. Similarly, another misleading spectral characteristic of the corona is caused by the layer’s great temperatures in its higher regions. That is, electrons in the ionized gas zip around with great velocity, and when they encounter and deflect photons the Doppler shift is so great that they produce a continuous spectrum (without absorption lines).
Gas in the lower corona is not as highly ionized as in the outer regions, though, and this tells astronomers that the gas is not as hot there. Just above the chromosphere, coronal gas has a temperature of about 500,000 k, but the temperature rises to over 2 x 106 k in the outermost regions. Coronagraph imagery shows magnetic streamers extending out into the corona from the photosphere, and astronomers believe that these and other of the Sun’s magnetic phenomena are responsible for making coronal and chromospheric gases so much hotter than photospheric gas. But, in order to understand why this is so, an explanation of the Sun’s magnetic activity is needed.
The Sun’s Magnetic Fields and Sunspots
No study of the Sun is complete without an in-depth look at its magnetic fields, as they play a crucial role in determining some of the star’s characteristic behaviors and weather. Also, the interactions between solar gases and magnetic forces are responsible for some pretty fantastic phenomena! While some aspects are still a mystery, astronomers have been able to make some important observations over the years that have led to plausible theories and models explaining the physics behind the Sun’s magnetic fields.
Recall from the discussion on nuclear fusion and the P-PC that much of the Sun’s gas is highly ionized. This is important because the great quantity of free electrons in ionized gases make them excellent conductors of electricity. Moreover, recall that energy flowing outward from the Sun must pass through the convective zone before escaping into space, and according to the laws of physics, high electrical conductivity and convection are two of the requisite criteria for creating electromagnetism. However, one more key component is needed: rapid rotation.
When these three conditions are met – i.e., electrical conductivity, convection and (rapid) rotation – some of the convective energy in a system may be converted into an electromagnetic field. This is known as the dynamo effect, and current theory holds that it is responsible for generating the Sun’s magnetic fields at the bottom of the convection zone. The situation is a bit more complex as far as the Sun is concerned though.
As mentioned earlier, the Sun’s photosphere has a rotational period of 25.38 days at the equator, but this statement is somewhat misleading because the Sun is a gaseous entity, and therefore it does not rotate uniformly as a rigid body would. For example, at 45 degrees latitude the photosphere has a rotational period of 27.8 days. Furthermore, observations in helioseismology have shown that interior regions of the Sun also rotate at different speeds. Physicists call this phenomenon differential rotation, and it has an important effect on the behavior of the Sun’s magnetic fields.
According to the Babcock model, differential rotation in the Sun serves to tangle the star’s magnetic fields in much the same way that loose string can get caught up in a spinning wheel. You can picture this if you imagine a single magnetic field on the Sun as a long piece of rope that is attached lengthwise to a rotating sphere. Only, the mid-section of the sphere rotates more quickly than its top or bottom sections. Consequently, the middle section of the rope makes a full rotation while the other sections lag behind. As the sphere continues to rotate, the rope becomes more and more tangled.
Moreover, rising and sinking currents of gas in the Sun’s convective zone serve to twist magnetic fields until they “kink up” like a dislodged bicycle chain. When this happens, the already tangled fields breach the photosphere in looped arches that are polarized at each end, much like a flat bar magnet that has been forced into a horseshoe shape and stood up. Because convection is disrupted where magnetic arches break through the Sun’s surface layer, the gas is slightly cooler in those areas, and we perceive them as dark regions on the photosphere, called sunspots.
In fact, sunspots are not really “cool” regions at all, as they have average temperatures of around 4200 k. They only appear dark amid the hotter regions of the photosphere because the amount of light radiated by a black body surface is directly proportional to its temperature raised to the fourth power (known as the Stefan–Boltzmann law). Still, astronomers use the same language to describe sunspots as other people do to describe shadows – i.e., the dark, innermost region of a sunspot is known as the umbra, and the comparatively light, outer region is known as the penumbra.
Sunspots occur frequently in pairs, which would seem to lend credence to the Babcock model if a pair represents the “legs” of an arched magnetic field. However, sunspots have also been observed in larger, more complex groupings. Curiously, the polarizations of sunspot groups are reversed north and south of the Sun’s equator, at any given time. For instance, if a pair of sunspots north of the equator are positive-negative, then another pair south of the equator will be negative-positive. Additionally, every eleven years the leading polarity of sunspot pairs is reversed (e.g., with negative-positive pairs north of the equator, and positive-negative pairs south of the equator).
Furthermore, at the beginning of any eleven year sunspot cycle, astronomers have observed that sunspots tend to appear less frequently and at higher solar latitudes. But in subsequent years sunspots begin to appear more frequently and closer to the equator. The Babcock model explains this as a consequence of differential rotation.
That is, at the beginning of a sunspot cycle the Sun’s magnetic fields are less tightly wound. Therefore, the regions of the fields that are twisted into protruding arches (i.e., sunspot pairs) by convection occur at higher latitudes, if at all. But as the fields become more tightly wound and tangled over time, arches protrude more frequently (causing over a hundred sunspots at maximum) and in tighter patterns.
Eventually, the magnetic fields become so chaotic that adjacent, oppositely polarized magnetic regions of the Sun begin to change themselves in order to sync with each other. When this happens, it is as though someone hits the “re-set” switch, and all of the Sun’s magnetic fields re-arrange themselves into simpler patterns and begin the cycle anew. Astronomers distinguish between an eleven year cycle for sunspots and a twenty-two year cycle for magnetic activity (i.e., two sunspot cycles), because it takes twenty-two years for the sunspot pairs north/south of the equator to return to their original leading polarities.
A common graph used by astronomers showing the relationship between sunspot latitudes over time is called the Maunder Butterfly diagram; so named because the groupings resemble butterfly wings when charted. From 1645 – 1715, during a period known as the Maunder minimum, very little sunspot activity was observed on the Sun, which corresponded with a period of unusually cold winters in Europe and North America (from around 1500 – 1850, known as the “little ice age”). This, and other similar observations, would suggest that there may be some connection between solar magnetic activity and the amount of solar energy that our planet receives.
Prominences, Flares, CMEs and the Solar Wind
Magnetic fields on the Sun are catalysts for some of the most exciting solar phenomena observable. For instance, the sunspots that our eyes perceive on the surface of the Sun are really just visual wavelength indicators that certain areas are exhibiting highly complex magnetic activity. This activity extends past the photosphere and chromosphere into the lower corona, and may even affect us here on Earth. Astronomers refer to the highly magnetic areas indicated by sunspots as active regions, and they may be observed more effectively at shorter, non-visual wavelengths (e.g., far-ultraviolet).
It is also possible to quantify the relative strengths of magnetic fields in and around sunspot regions by measuring the Zeeman effect, which is the degree to which the electrons in an atom are disrupted in the presence of a magnetic field, thereby allowing the atom to absorb photons of different wavelengths than would otherwise be possible – i.e., the more extreme the Zeeman effect, the stronger the magnetic field affecting the atom. Spectroscopic analysis of light emitted in sunspot regions shows how much the solar gases in those regions have been disrupted by magnetic activity.
Recall that the Sun’s magnetic fields are tangled by differential rotation and twisted by convection into arch-like protrusions. When ionized gases become trapped within these arches they are known as prominences(or filaments when seen from an overview angle against the surface of the photosphere), and they resemble great lassos of flame whipping around in the Sun’s atmosphere. Commonly extending as far out as the lower corona, prominences can be many times the size of Earth, and they last anywhere from a few hours (eruptive prominences) to days at a time (quiescent prominences). The gas in prominences may have temperatures of anywhere from 60,000 – 80,000 k, and though this is lower than the average temperatures in the chromosphere and corona, the whipping motion of prominences is thought to be responsible for heating the gases in these atmospheric layers to such high degrees.
In active regions, oppositely directed magnetic fields can meet and cancel each other out. This is known as a reconnection, and when the Sun’s magnetic fields neutralize each other in such a way, energy is released as high-energy particles (i.e., protons and electrons) and short-wavelength photons (i.e., ultraviolet and x-ray) in a solar flare. These phenomena resemble crashing tidal waves of flame that speed out from the surface of the Sun into space.
Photons from solar flares reach Earth at the speed of light (in a little over eight minutes), increasing the level of ionization in our planet’s atmosphere. This can sometimes disrupt communications that are carried on longer wavelength frequencies. For example, a common scene in science fiction movies involves someone listening to a radio that suddenly plays static for a few seconds before the signal clears up again. The character then rationalizes the temporary loss of reception as a consequence of solar flare activity, before continuing on about his or her business (right before the hostile aliens attack).
The high-energy protons and electrons from a flare can take as little as a few hours or as long as a few days to reach us. But when they do, our planet’s magnetic field can be distorted, satellites and electrical power lines can be damaged and even navigation systems might be impaired. The magnitude of a solar flare can vary, but in particularly severe cases the phenomena can be very disruptive to life on Earth.
On September 1, 1859, British astronomer Richard Carrington observed two brilliant points of light on a screen reflecting solar sunspot images. The points grew rapidly, taking on kidney-like shapes, and then receded after only five minutes. Carrington realized that he had observed some unique phenomenon, but he was not sure what to make of it.
Before dawn of the next day, telegraph systems around the world began malfunctioning. Electrical discharges from the machines shocked operators and started fires. Also, interconnecting networks of telegraph lines became so charged that they could transmit messages without a power source. Most surprisingly, the sky was lit up with brilliant auroras across the globe, and strange combinations of colors rippled through the atmosphere even at tropical latitudes. The extraordinary auroras were said to have been so bright that people could read newspapers by them as though it were the middle of the afternoon.
Lay people of the day must have been confounded by the spectacle, and no one could have known that the unusual events they were witnessing were caused by those mysterious points of light on the solar surface that had been observed by Carrington the previous day. In fact, we now know the flashes of light were the results of magnetic reconnections that triggered a massive solar flare – the largest to affect Earth in 500 years. High energy particles ejected from the flare reached Earth a day later, exciting the atoms and molecules of atmospheric gases into emitting energy as auroras, which in turn induced electric currents in telegraph lines.
Particularly violent reconnections, like the one that triggered the massive solar flare on September 1, 1859, release a great deal of energy. Sometimes, the energy is enough to blow a large amount of ionized gas from the corona on a collision course with Earth. These coronal mass ejections (or CMEs) slam into Earth’s magnetic field, distorting it and creating disruptions.
High energy particles and gas ejected from flares, as well as CMEs make up a part of the solar wind, but much of it comes from regions where the Sun’s magnetic field does not loop back into the photosphere. Rather, open-ended fields release currents of energy into the Sun’s atmosphere, pushing coronal gases out into space like giant, wind-generating turbines. Viewed at x-ray wavelengths, these regions of open-ended magnetic field activity on the Sun’s surface appear dark, and astronomers refer to them as “coronal holes.”
Our entire biosphere evolved under the nurturing care of the Sun. Every type of energy that is consumed here on Earth began as a gamma ray in that far-off, stellar core. It is small wonder that civilizations throughout human history have worshipped the star as a deity – i.e., the Sun literally makes life as we know it possible. Indeed, its energy serves as the metaphorical umbilical cord to our continued existence.
Similarly, with a total mass of 1.989 x 1030 kg’s and a circumference of about 4,373,034 km, the Sun seems incredibly vast by Terran standards (our planet’s mass and circumference are 5.9726 x 1024 kg and around 40,075 km, respectively).* In fact, it would take the equivalent of about 333,021 Earths to match the mass of the Sun:
(1.989 x 1030 kg)/(5.9726 x 1024 kg) = 333,020.795
Furthermore, the Sun converts roughly five million tons of matter into energy every second. But, even though this is more than enough to satisfy our needs here on Earth, the Sun will ultimately convert only about 0.07 percent of its total mass into energy before dying. This is largely due to the fact that helium nuclei are slowly accumulating in the star’s inner-core, as an unhappy consequence of the region’s radiative equilibrium (i.e., non-convective nature). Moreover, recall that the Sun is only able to sustain itself due to fusion reactions taking place inside its core, thanks to the heat and density conditions there (1.5 x 106 k and 160gm/c3, respectively).
However, the number of fusion reactions taking place per unit volume in the core is actually quite small. This might seem surprising for a region that holds forty percent of the Sun’s mass within only ten percent of its volume. Nevertheless, solar fusion reactions take place in only the hottest, densest regions of the stellar core – i.e., the center. As a consequence, the gathering helium nuclei will eventually saturate the inner-core like ash that has built up in a chimney, and P-PC reactions will no longer take place there. When this happens, the Sun will proceed to the red giant stage of its evolution.
Still, P-PC fusion occurs frequently enough within the core to produce the energy our Sun needs in order to resist being crushed by the force of its own gravity – i.e., Sol gets the job done. Though perhaps not the most streamlined model available, our stellar patron’s efforts are sufficient to power up its photosphere with enough luminosity (3.826 x 1026 J/s) to make life possible on Earth – not too shabby for an average star on the main sequence, with a spectral classification of G2 V! Moreover, there is no indication that the Sun will stop coming through for us any time soon, as scientists estimate that the star has another five billion years or so before it enters the red giant phase (give or take an epoch). So, while it is technically true that the Sun is not as efficient at converting matter into energy as some of its other celestial siblings (e.g., red dwarf stars), try not to judge it too harshly: like a single, middle-class parent who always finds a way to make ends meet, the Sun does the best that it can with what it has. Also, it’s the only star we’ve got!
* Earth’s radius at the Equator, or 6378.1 km, was used to derive our planet’s diameter and circumference measurements of 12,756.2 km and 40,075 km (i.e., d = 2r & c = 2p r), respectively. However, it should be noted that Earth is an oblate spheroid, not a perfect sphere. Therefore, the diameter and circumference measurements used here are only approximations.
** The binding energy released during the formation of a common helium nucleus is approximately 28.3 MeV, or 28.3(1.780 x 10-30 kg) = 5.0374 x 10-29 kg
*** Astronomers reason that such absorption lines must have come from the chromosphere, because it is unlikely that photons of those wavelengths would have been able to escape from the Sun’s lower layers.
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