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Our favorite star has to be the Sun. It is, by all accounts, just an average star, but extraordinary to us as a giver of life. This lesson covers two chapters in the text, starting first with the overall structure of the Sun and the Earth–Sun connection, and finishing with the nuclear powerhouse that rages in the core of the Sun. The activity connected to the lesson introduces the data we've obtained from the SOHO satellite, which has been monitoring the Sun continuously over the past few years. You will be calculating the tremendous velocity of a solar flare and the time it will take for the energy released to reach and affect Earth.
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After completing this lesson (including the lecture), you should be able to
Solar Interior
This lesson takes a look at the Sun at two different levels. The first level considers how the Sun affects the Earth and the life on it—the activity of the Sun and its 22-year sunspot cycle. Solar activity involves the photosphere, chromosphere, and corona of the Sun. For the second level, we are going to imagine that we have an indestructible spaceship that is capable of withstanding temperatures of millions of degrees and extremely high pressure. We are going to imagine that we are able to journey not only to the Sun but also to its very core, right to where the nuclear powerhouse is raging. Along the way, we'll take notes of what we see out our window.
Read through Sections 6.2 through 6.5 carefully to get a sense of what all solar activity involves, noting in particular the "Making Connections" on page 148. Solar activity is not just the presence of sunspots on the surface of the Sun. It includes these sunspots, of course, but also includes plages, prominences, and flares. It is these flares that, if aimed towards the Earth, can wreak havoc with our communication system, our power supplies, and can be potentially lethal to our astronauts aboard a space shuttle or the international space station.
Include in your reading about the Sun–Earth connection the information found at " A Primer on Space Weather," from NOAA (National Oceanic and Atmospheric Administration) Space Environment Center. The site starts with a brief summary about the Sun and then goes into more details about aurora, proton events, geomagnetic storms, and how various systems on Earth can be disrupted. Obviously, we cannot control the Sun's activity. The goal of scientists now is to gain enough knowledge about the active Sun to be able to predict solar flares and how to better safeguard our systems from the effects of a "supershot" of radiation. As the Space Weather Advisories homepage states: "Space Weather Advisory Alerts and Warnings are being developed. Alerts will be issued when a significant space weather event has been observed, is in progress, or has ended; Warnings will be issued when a significant space weather event is imminent, or likely. " This is about the best we can do right now.
As we leave our planet Earth and head towards the Sun (also affectionately known as "Sol"), our star appears as a brilliant orb against a now black sky. It is because the Earth's atmosphere scatters that Sun's light that we have a bright sky during the day and cannot see any other stars. Take a look at the image at the right taken by the SOHO satellite, and you'll see that many hundreds of stars appear around the Sun. The bright star, Antares, and the planet Mercury are identified in the image. The full light of the Sun is blocked by a round disk in front of the optics of the satellite so that the light detectors are not damaged. (Image: NASA/SOHO)
The Sun's corona is seen streaming away from the Sun in this image. The corona may extend many millions of kilometers away from the Sun. When the Sun is around its peak in activity, the corona forms these extensive streamers. The corona diminishes significantly during the inactive part of the Sun's activity cycle. The corona is extremely hot: around 1-2 million degrees Kelvin. This temperature is much, much greater than that of the Sun's photosphere, or what we call its surface. How can that be? Where is the heat coming from? Recent observations by NASA's Transition Region and Coronal Explorer (TRACE) spacecraft indicate that the heating takes place at the base of millions of coronal loops. The strong magnetic field present in these loops may play a significant role in the heating.
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Even though the temperature of the corona is extreme (almost as hot as the core of the Sun!), its heat content is very low. This is because its density is about one-billionth of that of the Earth's atmosphere at sea level. Thus, we will be able to pass through the corona without worrying about getting scorched. Comparing the core of the Sun (where we are more likely to get scorched) with its corona is similar to comparing a pot of boiling water (100 degrees Celsius, or 212 degrees Fahrenheit) to an oven of the same temperature. Where you would never put your hand in the pot, you would have little hesitation sticking it in the middle of the oven. Heat content depends not only on the temperature, but also the density of the material.
Why don't we see the corona on a normal sunny day? This is because it is so diffuse that the bright light of the Sun overwhelms it. If we could view the Sun in x-rays from above the atmosphere, we would notice that it looks patchy, with bright spots and dark areas (see Fig. 6.10 of the text). The dark areas are called coronal holes, and are the regions where the Sun expels particles that make up the solar wind. In the activity connected with this lesson, you will be measuring the velocity of a coronal mass ejection, material that is escaping from one of these holes and doing so quite violently. It is the particles contained in the solar wind—electrons, protons, helium nuclei—that interact with the magnetic field of the Earth and create the spectacular northern and southern lights.
A
coronal mass ejection. The bright object to the lower right of the Sun is the
planet Venus—so bright that it floods the light detector and causes the
long, narrow "bar." The circular asterism at the bottom is the handle
of the teapot—Sagittarius. (Image: NASA/SOHO)
We have survived the trip through the corona; there were no coronal mass ejections pointed our way to inundate our electronics with charged particles. Our next region is the chromosphere.
The chromosphere lies above the photosphere. The temperature ranges from roughly 10,000 K to 500,000 K or more. Since the chromosphere represents a hot diffuse gas, astronomers will detect many emission lines from the elements present there. (Once the light from the Sun is blocked, astronomers may detect many emission lines when observing the corona, some representing elements that have lost 3, 4, 5, or even more electrons. Hot indeed!)
Four features really characterize the chromosphere: solar flares, prominences, plages, and spicules.
An interesting note about the chromosphere. The chromosphere, being hotter than the photosphere below it, shows emission spectra representing the various elements that make up the Sun. In the year 1868, while spectroscopy of the Sun was still in its infancy, astronomers noted an unusual emission line while observing a prominence. This line had not been seen on Earth, and could not be reproduced, at that time, in the laboratory. A British civil servant, Norman Lockyer, suggested that it was a new element. The element was, of course, helium, for helios meaning Sun. The element was later found on Earth in 1895. A strange tale of discovery for the second most common element in the Universe.
At last we have reached what we call the "surface" of the Sun: the photosphere or sphere of light. The photosphere defines the solar disc and gives the Sun its characteristic temperature and color: 5800 Kelvin and green-yellow. (Our atmosphere scatters some of the blue light from the radiation coming from the Sun, making the Sun look more yellow from the ground.)
The photosphere has a number of notable features: granulation and sunspots being the main ones.
If we assume that the sunspots are, in fact, embedded in the photosphere of the Sun, we could, theoretically, determine the rotation period of the Sun. This is exactly what scientists have done. We find out a curious thing: the Sun rotates faster at its equater than at its poles. It takes about 25 days for a complete rotation at the equator, while it takes closer to 30 days at the poles. This is called differential rotation. The Earth and any other solid body cannot rotate freely like this, only gaseous objects can (like stars and the gas giants of the solar system). It's this differential rotation that causes the magnetic field lines of the Sun to get all wrapped up and eventually pop out.
We are going to enter the photosphere through one of the sinking edges of a granule, on our way to the convection zone, but first let's take a closer look at granulation and some sunspots in these images from the Big Bear Solar Observatory. Click on each image for a closer look.
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| This image is taken at near infrared wavelengths. Each cell is about 1000 kilometers across. March 12, 1999 | A close-up image of a highly active sunspot taken in white light (no filter on the telescope). We see clear down to the granular pattern in the photosphere. June 14, 1989 | This high resolution image was taken by isolating the 656.3 nm wavelength of neutral hydrogen. Follow the arches from the spot to plage above. July 4, 1974 |
Of the three means of transporting radiation—convective, radiative, and conductive—convective transport is the most efficient. When there is a large temperature gradient—meaning when the temperature difference between the bottom of a layer and the top of the layer is very great—it takes convective action to transport the hot gas to the next layer for cooling. We are all familiar with convection when we boil water in a pan. Stars like the Sun have a convective layer near their surfaces. Massive stars, roughly ten times or more massive than the Sun, have convective layers outside their cores because their nuclear fusion is producing a huge amount of energy that needs to get out. No mass is gained or loss in regions where convection is occurring, just the radiation is passed to the next layer. Our theories and models of how convection works are inadequate at this time, even though a number of astronomers have been working on the problem. The models most astronomers work with are nearly 50 years old; this is indeed outdated as far as theoretical astrophysics is concerned.
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Here is a sketch of our spacecraft diving into the convective zone.
We will have to be quick once we get to the bottom of the convective zone not to be carried right back up to the photosphere! |
If the layer does not have a high temperature difference between the bottom and the top, the only means to get the heat out is through radiative transfer (a complicated process indeed if we were to go into it in greater detail). Conductive transports needs to have a high density material, high enough for atoms and molecules to bump against each other. As the text mentions, radiation is not an efficient means of transporting energy in stars because the photons can only go a tiny distance, around 1 cm or so, before being reabsorbed by another atom and emitted in a random direction. Each time that this process happens, the photon that is remitted may or may not be of the same energy as the photon that was absorbed. Recall from the chapter on radiation that an electron excited to an upper level can cascade down levels, emitting lower-energy photons than the original exciting one. This ensures that the gamma rays produced by fusion in the core leave the Sun mostly as visible light about 1 million years later.
View a short animation of a photon in its random walk.
The process by which the photons zigzag their way through the radiative zone is called a random walk. Imagine that you are at a large concert, in the center of a large crowd, with a number of balloons being bounced around. If each participant hits the balloons in a random direction each time, how long would it take for any one balloon to go from the center to the edge (and over the wall)?
It is here, in the heart of the Sun, maybe 10% of its total volume, where the energy that supports all of life on Earth is produced. Thermonuclear fusion supplies the luminosity of the Sun, and has been consistently doing so on a steady basis for about 5 billion years.
How does this work? Your textbook provides a good explanation, bringing in the simple equation that unlocked the secret of the Sun: E = mc2. On its fundamental level, this equation means that energy and mass are interchangeable (that's what the equal sign means), to within a factor of c2. An unbelievably huge amount of energy can be obtained from very little mass. (On the other hand, it takes an enormous amount of energy, such as was available shortly after the Big Bang, to make matter out of energy.)
The mass lost during a complete cycle of the proton-proton cycle equals 0.048 × 10-27 kg. Take that change in mass times c2 and you get about 4.3 × 10-12 Joules per second. (One Joule per second equals one Watt.) The luminosity of the Sun, the amount of power it produces, is about 4 × 1026 Watts. Therefore, the Sun must be experiencing about 1038 conversions per second in its core, or the conversion of 2 × 109 kg mass into energy per second. It's a good thing the Sun has lots and lots of mass!
Now, I know you all are thinking: What if an "average" student were to suddenly have all or her or his mass converted to energy within 1 second?
You can calculate the amount of energy very easily. One kilogram is equal to 2.2 lbs, so that a 150-pound person has a mass of about 70 kg. Multiply this by the speed of light squared: 9 × 1016 and you get 6.3 × 1018 Joules. Since it happened in one second, that equals 6.3 × 1018 watts. Let's say the average light bulb gives off 60 watts. So, the average student could power about 1017 light bulbs for one second (each person on Earth would have a hundred-million light bulbs lit for one second), or a light bulb for one person on Earth for about 3 billion years, or about 10 billion light bulbs for one year.
The demonstration below shows the annihilation of the positron by an electron almost immediately after the positron is produced in the initial step. This annihilation of antimatter and matter produces energy; E = mc2 again describes where the energy comes from as the mass from the positron and electron is converted to energy.
View a simple animation of the proton-proton cycle.
The Sun is opaque below the photosphere. This means that we cannot see very deep into the Sun, so how do we know all of this? (Be sure to read your text carefully in the section where opacity is discussed.)
Theoretical stellar astrophysics uses mathematical equations and solutions to those equations at each layer of the Sun's interior in an attempt to match theory with the observations. The calculations keep very fast computers busy for days. At each layer of the Sun, the theorist works with the temperature, pressure, opacity (how easily can the photons travel through the layer), and radiative flux or convective flux. (Flux is the amount of radiation passing through each square cm of a layer per second.) Recently, previously classified calculations of nuclear cross sections (that is, how easily the elements will fuse) done at Los Alamos and Jet Propulsion Laboratory were released. The theory took a giant leap forward.
Understanding what the "standard solar model" is and how it is calculated is a bit on the complicated side, but not impossible to understand. For an extremely readable discussion checkout The Standard Solar Model by D. B. Guenther and P. Demarque, two of the world's experts in the field of stellar astrophysics.
From the textbook:
Sun-Earth Connection:
Chromosphere:
Photosphere:
Solar Theory:
Solar Observatories:
