Embed Size (px)
Citation preview
1
UFABC – NHZ3043 – NOÇÕES DE ASTRONOMIA E COSMOLOGIA – Curso 2016.2 Prof. Germán Lugones
Capítulo 8 Meio interestelar
Nebulosa Cabeça do Cavalo
2
‣A pesar de que a maior parte da massa da Via Láctea está contida nas estrelas, o meio interestelar não é completamente vazio. Principalmente no disco da Galáxia, o meio interestelar contém gás e poeira, distribuídos na forma de nuvens individuais, e também em um meio difuso.
‣A densidade típica do meio interestelar é de um átomo de hidrogênio por cm3 e, aproximadamente, 100 grãos de poeira por km3.
‣O gás interestelar constitui, aproximadamente, 10% da massa da Via Láctea, ao passo que a poeira agrupa menos de 1% da massa em gás.
‣Raios cósmicos, que são partículas altamente energéticas, estão misturados com o gás e a poeira, e existe ainda um campo magnético galático, fraco (≃ 10 μG).
‣Atualmente, as observações mais importantes do meio interestelar são feitas em comprimentos de onda de rádio e de infravermelho, uma vez que o pico de emissão encontra-se frequentemente nestes comprimentos de onda.
3
450 CHAPTER 18 The Interstellar Medium
The matter among the stars is collectively termed the interstellar medium. It is made up of two components—gas and dust—intermixed throughout all space. The gas is made up mainly of individual atoms, of average size 10−10 m (0.1 nm) or so, and small molecules, no larger than about 10−9 m across. Interstellar dust is more complex, consisting of clumps of atoms and molecules—not unlike chalk dust or the microscopic particles that make up smoke, soot, or fog.
Apart from numerous narrow atomic and molecular absorption lines, the gas alone does not block radiation to any great extent. The obscuration that is evident in Figure 18.1 is caused by the dust. Light from distant stars cannot penetrate the densest accumulations of interstellar dust any more than a car’s headlights can illuminate the road ahead in a thick fog.
Extinction and ReddeningWe can use its effect on starlight to measure both the amount and the size of interstellar dust. As a rule of thumb, a beam of light can be absorbed or scattered only by particles hav-ing diameters comparable to or larger than the wavelength of the radiation involved. Thus, a range of dust particle sizes will tend to block shorter wavelengths most effectively. Furthermore, even for particles of a given size, the amount of obscuration (that is, absorption or scattering) produced by particles of a given size increases with decreasing wave-length. The size of a typical interstellar dust particle—or dust grain—is about 10−7 m (0.1 μm), comparable in size to the wavelength of visible light. Consequently, dusty regions of interstellar space are transparent to long- wavelength radio and infrared radiation, but opaque to shorter wave-length optical and ultraviolet radiation. The overall dim-ming of starlight by interstellar matter is called extinction.
18.1 Interstellar MatterFigure 18.1 is a mosaic of photographs covering a much greater expanse of universal “real estate” than anything we have studied thus far. From our vantage point on Earth, the panoramic view shown here stretches all the way across the sky. On a clear night, it is visible to the naked eye as the Milky Way. In Chapter 23, we will come to recognize this band as the flattened disk, or plane, of our Galaxy.
The bright regions in this image are congregations of innumerable unresolved stars, merging together into a con-tinuous blur at the resolution of the telescope. However, the dark areas are not simply “holes” in the stellar distribution. They are regions of space where interstellar matter obscures (blocks) the light from stars beyond, blocking from our view what would otherwise be a rather smooth distribution of bright starlight. Their very darkness means that they cannot easily be studied by the optical methods used to examine stellar matter. There is, quite simply, nothing to see!
Gas and DustFrom Figure 18.1 (see also Figure 18.4), it is evident that interstellar matter is distributed very unevenly throughout space. In some directions, the obscuring matter is largely absent, allowing astronomers to study objects literally bil-lions of parsecs from the Sun. In other directions, there are small amounts of interstellar matter, so the obscuration is moderate, preventing us from seeing objects more than a few thousand parsecs away, but still allowing us to study nearby stars. Still other regions are so heavily obscured that starlight from even relatively nearby stars is completely absorbed before reaching Earth.
▲ FIGURE 18.1 Milky Way Mosaic The Milky Way Galaxy photographed panoramically, across 360° of the entire southern and northern celestial sphere. This band, which constitutes the central plane of our Galaxy, contains high concentrations of stars, as well as interstellar gas and dust. The white box shows the field of view of Figure 18.4. (ESO/S. Brunier)
R I U X GV
‣O gás está fortemente concentrado no plano galáctico e nos braços espirais. Assim, nessas regiões há muitos lugares onde as quantidade de massa nas estrelas e na matéria interestelar são aproximadamente iguais.‣ Figura: panorâmica da Via Láctea (360∘) mostrando estrelas, gás e poeira.
SECTION 18.2 Emission Nebulae 453
18.2 Emission NebulaeFigure 18.4 shows a magnified view of the central part of Figure 18.1 (the region indicated by the rectangle in the earlier figure), in the general direction of the constellation Sagittarius. The field of view is mottled with stars and interstellar matter. The patchiness of the obscuration is evident. In addition, several large fuzzy patches of light are clearly visible. These fuzzy objects, labeled M8, M16, M17, and M20, correspond to the 8th, 16th, 17th, and 20th objects in a catalog compiled by Charles Messier, an 18th-century French astronomer.* Today they are known as emission nebulae—glowing clouds of hot interstellar matter. Figure 18.5 enlarges the left side of Figure 18.4, showing the nebulae more clearly.
Observations of Emission NebulaeHistorically, astronomers have used the term nebula to refer to any “fuzzy” patch (bright or dark) on the sky—any region of space that was clearly distinguishable through a telescope, but not sharply defined, unlike a star or a planet. We now know that many (although not all) nebulae are clouds of interstellar dust and gas.
If a cloud happens to obscure stars lying behind it, we see it as a dark patch on a bright background, as in Figures 18.1, 18.2(b), and 18.4—a dark nebula. But if something within the cloud—a group of hot young stars, for example— causes it to glow, then we see a bright emission nebula
Astronomers infer this elongated structure from the fact that the light emitted by stars is dimmed and partially polarized, or aligned, by the intervening dust. Recall from Chapter 3 that light consists of electromagnetic waves com-posed of vibrating electric and magnetic fields. (Sec. 3.2, Fig. 3.7) Normally, these waves are randomly oriented, and the radiation is said to be unpolarized. Stars emit unpolar-ized radiation from their photospheres. However, under the right conditions, the radiation can become polarized en route to Earth, with the electric fields all vibrating in roughly the same plane. One way in which this can happen is if the radiation interacts with an elongated dust grain, which tends to absorb electric waves vibrating parallel to its length.
Thus, if the light detected by our telescope is polarized, it is because some interstellar dust lies between the emitting object and Earth. Based on this reasoning, astronomers have determined not only that interstellar dust particles must be elongated in shape, but also that they tend to be aligned over large regions of space.
The alignment of the interstellar dust is the subject of ongoing research among astronomers. The current view, accepted by most, is that the dust particles are affected by a weak interstellar magnetic field, perhaps a million times weaker than Earth’s field. Each dust particle responds to the field in much the same way that small iron filings are aligned by an ordinary bar magnet. Measurements of the blockage and polarization of starlight thus yield information about the size and shape of interstellar dust particles, as well as about magnetic fields in interstellar space.
CONCEPT Check
4 If space is a near-perfect vacuum, how can there be enough dust in it to block starlight?
*Messier was actually more concerned with making a list of celestial ob-jects that might be confused with comets, his main astronomical interest. However, the catalog of 109 “Messier objects” is now regarded as a much more important contribution to astronomy than any comets Messier dis-covered.
M20
M8
R I V U X G
▶ FIGURE 18.4 Milky Way in Sagittarius Enlargement of the central part of Figure 18.1, showing regions of brightness (vast fields of stars) as well as regions of darkness (where interstellar matter obscures the light from more distant stars). The field of view is about 35° across. Two of the emission nebulae discussed in the text are labeled. (ESO/S. Guisard)
AN
IMAT
ION
/VID
EO O
rion
Neb
ula
Mos
aic
4
‣ O gás interestelar é constituído, na maior parte, por hidrogênio neutro, que não é luminoso. Também pode haver alguns átomos e moléculas. ‣ As regiões com gás são transparentes a todos os comprimentos de onda com exceção de linhas de absorção estreitas, i.e. o gás praticamente não bloqueia a radiação.‣ Porém, perto de estrelas muito quentes e massivas, o hidrogênio é ionizado pela radiação ultravioleta provinda das estrelas e brilha por fluorescência. ‣Se existe suficiente hidrogênio ao redor dessas estrelas, ele será visível como uma nebulosa gasosa de emissão, brilhante, chamada região HII, ou nebulosa de emissão. ‣Um exemplo desse tipo de nebulosa é a Nebulosa de Órion, que se encontra a 1500 anos-luz da Terra.
Gás interestelar
5
458 CHAPTER 18 The Interstellar Medium
very different from conditions on Earth and warn us of the potential difficulties involved in extending our ter-restrial experience from our laboratories to the study of interstellar space.
Some regions of interstellar space contain extremely dilute, even hotter gas than is found within emission nebu-lae. Ultraviolet observations by space-based instruments have found that these superheated interstellar “bubbles,” making up the intercloud medium, may extend far into interstellar space beyond our local neighborhood and, con-ceivably, into the even vaster spaces among the galaxies. This high-temperature gas is probably the result of the violent expansion of debris from stars that exploded long ago. Somewhat like the Sun’s faint corona, these regions are dark despite their high temperatures because the density of matter there is very low. (Sec. 16.3)
The Sun seems to reside in one such low-density region—a huge cavity called the “Local Bubble,” sketched in Figure 18.11. The Local Bubble contains about 200,000 stars and extends for nearly 100 pc. It was probably carved out by multiple supernova explosions (see Chapters 20 and 21) that occurred several hundred thousand years ago in the Scorpius–Centaurus association, a rich cluster of bright young stars. Perhaps our hominid ancestors may have seen these ancient events—stellar catastrophes as bright as the full Moon—that now aid modern astronomers.
R I V U X G(b) (c)(a)
1 light-year
▲ FIGURE 18.10 Orion Nebula (a) Lying some 1400 light-years from Earth, the Orion Nebula (M42) is visible to the naked eye as the fuzzy middle “star” of Orion’s sword. (b) Like all emission nebulae, the Orion Nebula consists of hot, glowing gas powered by a group of bright stars in the center. In addition to exhibiting red H! emission, parts of the nebula show a slight greenish tint, caused by a so-called forbidden transition in ionized oxygen. (c) A high-resolution image shows rich detail in a region about 0.5 light-year across. Structural details are visible down to a level of 0.1”, or 6 light-hours—a scale comparable to the dimensions of our solar system. (NASA; ESO)
Hyadesstar cluster
Sirius
Aldebaran
SunVega
To Galacticcenter
Procyona Centauri
Arcturus
Scorpius-Centaurusassociation
130 light-years
▲ FIGURE 18.11 Local Bubble The Sun resides in a vast low-density region of space that engulfs us nearly spherically. This cavity was likely caused by stellar explosions long ago, which then heated the nearby interstellar gas and expelled it well out of the solar neighborhood. Several prominent stars in our nighttime sky are plotted in this artist’s conception, which depicts what the “bubble” might look like from afar.
(a) Lying some 1400 light-years from Earth, the Orion Nebula (M42) is visible to the naked eye as the fuzzy middle “star” of Orion’s sword.
(b) Like all emission nebulae, the Orion Nebula consists of hot, glowing gas powered by a group of bright stars in the center. In addition to exhibiting red H emission, parts of the nebula show a slight greenish tint, caused by a so-called forbidden transition in ionized oxygen.
(c) A high-resolution image shows rich detail in a region about 0.5 light-year across. Structural details are visible down to a level of 0.1”, or 6 light-hours—a scale comparable to the dimensions of our solar system.
6
320
15. The Interstellar Medium
Table 15.3. Element abundances in the interstellar mediumtowards ζ Ophiuchi and in the Sun. The abundances are givenrelative to that of hydrogen, which has been defined to be
1,000,000. An asterisk (*) means that the abundance has beendetermined from meteorites. The last column gives the ratioof the abundances in the interstellar medium and in the Sun
Atomic Name Chemical Interstellar Solar Abundancenumber symbol abundance abundance ratio
1 Hydrogen H 1,000,000 1,000,000 1.002 Helium He 85,000 85,000 ≈ 13 Lithium Li 0.000051 0.00158* 0.0344 Beryllium Be < 0.000070 0.000012 < 5.85 Boron B 0.000074 0.0046* 0.0166 Carbon C 74 370 0.207 Nitrogen N 21 110 0.198 Oxygen O 172 660 0.269 Fluorine F – 0.040 –
10 Neon Ne – 83 –11 Sodium Na 0.22 1.7 0.1312 Magnesium Mg 1.05 35 0.03013 Aluminium Al 0.0013 2.5 0.0005214 Silicon Si 0.81 35 0.02315 Phosphorus P 0.021 0.27 0.07916 Sulfur S 8.2 16 0.5117 Chlorine Cl 0.099 0.45 0.2218 Argon Ar 0.86 4.5 0.1919 Potassium K 0.010 0.11 0.09420 Calcium Ca 0.00046 2.1 0.0002221 Scandium Sc – 0.0017 –22 Titanium Ti 0.00018 0.055 0.003223 Vanadium V < 0.0032 0.013 < 0.2524 Chromium Cr < 0.002 0.50 < 0.00425 Manganese Mn 0.014 0.26 0.05526 Iron Fe 0.28 25 0.01127 Cobalt Co < 0.19 0.032 < 5.828 Nickel Ni 0.0065 1.3 0.005029 Copper Cu 0.00064 0.028 0.02330 Zinc Zn 0.014 0.026 0.53
elements. This interpretation is supported by the obser-vation that in regions where the amount of dust is smallerthan usual, the element abundances in the gas are closerto normal.
Atomic Hydrogen. Ultraviolet observations have pro-vided an excellent way of studying interstellar neutralhydrogen. The strongest interstellar absorption line, ashas already been mentioned, is the hydrogen Lymanα line (Fig. 15.15). This line corresponds to the transi-tion of the electron in the hydrogen atom from a statewith principal quantum number n = 1 to one with n = 2.The conditions in interstellar space are such that almostall hydrogen atoms are in the ground state with n = 1.Therefore the Lyman α line is a strong absorption line,whereas the Balmer absorption lines, which arise fromthe excited initial state n = 2, are unobservable. (The
Balmer lines are strong in stellar atmospheres with tem-peratures of about 10,000 K, where a large number ofatoms are in the first excited state.)
The first observations of the interstellar Lyman α linewere made from a rocket already in 1967. More exten-sive observations comprising 95 stars were obtained bythe OAO 2 satellite. The distances of the observed starsare between 100 and 1000 parsecs.
Comparison of the Lyman α observations with ob-servations of the 21 cm neutral hydrogen line have beenespecially useful. The distribution of neutral hydrogenover the whole sky has been mapped by means of the21 cm line. However, the distances to nearby hydrogenclouds are difficult to determine from these observa-tions. In the Lyman α observations one usually knowsthe distance to the star in front of which the absorbingclouds must lie.
The abundances are given relative to that of hydrogen, which has been defined to be 1,000,000. The last column gives the ratio of the abundances in the interstellar medium and in the Sun.
7
‣ A poeira interestelar tem uma composição mais complexa. Ela é feita de aglomerados de átomos e moléculas, de forma similar à fumaça, névoa, poeira de giz em grãos de vários tamanhos, mas muito menores (~1µm) do que a poeira aqui na Terra. Ela é contém principalmente grafite, silicatos e gelo de água. ‣ A luz das estrelas mais distantes não penetra a poeira densa, da mesma forma que um farol de carro não penetra neblina densa. ‣ A poeira circundando estrelas reflete a luz formando uma nebulosa de reflexão, de cor azulada. O espectro dessas nebulosas é o mesmo da estrela que a ilumina.
Poeira interestelar
452 CHAPTER 18 The Interstellar Medium
atomic or molecular hydrogen; some 9 percent is helium, and the remaining 1 percent consists of heavier elements. The abundances of some heavy elements, such as carbon, oxygen, silicon, magnesium, and iron, are much lower in interstellar gas than in our solar system or in stars. The most likely expla-nation for this finding is that substantial quantities of these elements have been used to form the interstellar dust, tak-ing them out of the gas and locking them up in a form that is much harder to observe.
In contrast to interstellar gas, the composition of inter-stellar dust is currently not very well known. We have some infrared evidence for silicates, graphite, and iron—the same elements that are underabundant in the gas—lending sup-port to the theory that interstellar dust forms out of inter-stellar gas. The dust probably also contains some “dirty ice,” a frozen mixture of ordinary water ice contaminated with trace amounts of ammonia, methane, and other chemical compounds. This composition is quite similar to that of cometary nuclei in our own solar system. (Sec. 14.2)
Dust ShapeCuriously, astronomers know the shapes of interstellar dust particles better than their composition. Although the minute atoms in the interstellar gas are basically spheri-cal, the dust particles are not. Individual dust grains are apparently elongated or rodlike, as shown in Figure 18.3(a), although recent theoretical studies of how dust particles col-lide, stick, and break up suggest that their larger scale struc-ture may be considerably more complex (Figure 18.3b).
Overall DensityGas and dust are found everywhere in interstellar space—no part of our Galaxy is truly devoid of matter. However, the den-sity of the interstellar medium is extremely low. Overall, the gas averages roughly 106 atoms per cubic meter—just 1 atom per cubic centimeter—although there are large variations from place to place: Densities ranging from 104 to 109 atoms/m3 have been found. Matter this diffuse is far less dense than the best vacuum—about 1010 molecules/m3—ever attained in laboratories on Earth.
Interstellar dust is even rarer. On average, there is only one dust particle for every trillion or so atoms—just 10−6 dust particles per cubic meter, or 1000 per cubic kilometer. Some parts of interstellar space are so thinly populated that harvesting all the gas and dust in a region the size of Earth would yield barely enough matter to make a pair of dice.
How can such fantastically sparse matter diminish light radiation so effectively? The key is size—interstellar space is vast. The typical distance between stars (1 pc or so in the vicin-ity of the Sun) is much, much greater than the typical size of the stars themselves (around 10−7 pc). Stellar and planetary sizes pale in comparison to the vastness of interstellar space. Thus, matter can accumulate, regardless of how thinly it is spread. For example, an imaginary cylinder 1 m2 in cross section and extending from Earth to Alpha Centauri would contain more than 10 billion billion dust particles. (Sec. 17.1) Over huge distances, dust particles accumulate slowly, but surely, to the point at which they can effectively block visible light and other short-wavelength radiation. Even though the den-sity of matter is very low, interstellar space in the vicinity of the Sun contains about as much mass as exists in the form of stars.
Despite their rarity, dust particles make interstellar space a relatively dirty place. Earth’s atmosphere, by comparison, is about a million times cleaner. Our air is tainted by only one dust particle for about every billion billion (1018) atoms of atmospheric gas. If we could compress a typi-cal parcel of interstellar space to equal the density of air on Earth, this parcel would contain enough dust to make a fog so thick that we would be unable to see our hand held at arm’s length in front of us.
CompositionThe composition of interstellar gas is reasonably well understood from spectroscopic studies of absorption lines formed when light from a dis-tant star interacts with gas along the observer’s line of sight (see Section 18.3). In most cases, the elemental abundances detected in interstellar gas mirror those found in other astronomical objects, such as the Sun, the stars, and the jovian planets. Most of the gas—about 90 percent by number—is
10–7 m(a) (b)
Grains are linear, orrodlike, on smallscales, c
cbut can becometangled and twistedin complex ways onlarger scales.
▲ FIGURE 18.3 Interstellar Dust (a) Typical interstellar dust particles, as inferred from polarization studies, have sizes of only about one ten-thousandth of a millimeter, yet space contains enough of them to obscure our view in certain directions. (b) This result of a computer simulation shows how grains may grow as dust particles collide, stick, and fragment in interstellar space.
8
‣As partículas de poeira, com tamanhos de 0,1 a 1 mícron, são suficientemente pequenas para espalharem (desviar a direção, sem absorver) a luz de menor comprimento de onda (luz azul, λ ≤ 0,4 µm) mais eficientemente do que as de maior comprimento de onda (luz vermelha, λ ≥ 0,7 µm). ‣De fato, fótons azuis são desviados cerca de 10 vezes mais eficientemente do que os fótons vermelhos. Quando um fóton é desviado, sua direção muda aleatoriamente.
SECTION 18.2 Emission Nebulae 455
Figure 18.9 shows enlargements of two of the nebu-lae visible in Figure 18.5. Notice again the hot, bright stars embedded within the glowing nebular gas and the pre-dominant red coloration of the emitted radiation in parts (a) and (c). The relationship between the nebulae and their dust lanes is again evident in Figures 18.9(b) and (d), where regions of gas and dust are simultaneously silhouetted against background nebular emission and illuminated by foreground nebular stars.
The interaction between stars and gas is particularly striking in Figure 18.9(b). The three dark “pillars” vis-ible in this spectacular Hubble Space Telescope image are part of the interstellar cloud from which the stars formed. The rest of the cloud in the vicinity of the new stars has already been heated and dispersed by their radiation in a process known as photoevaporation. The fuzz around the edges of the pillars, especially at the top right and center, is the result of this ongoing process. (See also an up-close view of another such pillar in M16 in the chapter-opening photo.) As photoevaporation continues, it eats away the less dense material first, leaving behind delicate sculp-tures composed of the denser parts of the original cloud, just as wind and water create spectacular structures in Earth’s deserts and shores by eroding away the softest rock. The process is a dynamic one: The pillars will even-tually be destroyed, but probably not for another hundred thousand or so years.
Spectroscopists often refer to the ionization state of an atom by attaching a roman numeral to the chemical symbol for the atom—I for the neutral (that is, not ionized) atom, II for a singly ionized atom (an atom missing one electron), III
producing huge amounts of ultraviolet light. As ultra-violet photons travel outward from the star, they ionize the surrounding gas. As electrons recombine with nuclei, they emit visible radiation, causing the gas to f luoresce, or glow. (Sec. 4.2) The reddish hue of these nebu-lae—and, in fact, of all emission nebulae—results when hydrogen atoms emit light in the red part of the visible spectrum. Specifically, it is caused by the emission of radiation at 656.3 nm—the Hα line discussed in Chap-ter 4. (More Precisely 4-1) Other elements in the nebula also emit radiation as their electrons recombine, but because hydrogen is so plentiful, its emission usually dominates.
Woven through the glowing nebular gas, and plainly visible in Figures 18.5–18.7, are lanes of dark, obscuring dust. These dust lanes are part of the nebulae and are not just unrelated dust clouds that happen to lie along our line of sight. The bluish region visible in Figures 18.6 and 18.7 immediately above M20 is another type of neb-ula unrelated to the red emission nebula itself. Called a reflection nebula, it is caused by starlight scattered from dust particles in interstellar clouds located just off the line of sight between Earth and the bright stars within M20. Reflection nebulae appear blue for much the same reason that Earth’s daytime sky is blue: short-wavelength blue light is more easily scattered by interstellar matter back toward Earth and into our detectors. (More Pre-cisely 7-1) Figure 18.8 sketches some of the key features of emission nebulae, illustrating the connection between the central stars, the nebula itself, and the surrounding interstellar medium.
Light scattered through adusty cloud, not along the
line of sight, can look bluer.
Red light is emitted bynebulae when electrons
and protons recombine toform hydrogen atoms.
Hotstar(s)
Dust lane
Ultravioletradiation
Darkinterstellar
cloud
Visiblestarlight
Re-emittedvisible light
EMISSIONNEBULA
Dusty cloud
Ionizedgas
Observer
Unscatteredred light
REFLECTIONNEBULAScattered
blue light
▲ FIGURE 18.8 Nebular Structure An emission nebula results when ultraviolet radiation from one or more hot stars ionizes part of an interstellar cloud. If starlight happens to encounter another dusty cloud, some of the radiation, particularly at the shorter wavelength blue end of the spectrum, may be scattered back toward Earth, forming a reflection nebula.
9
460 CHAPTER 18 The Interstellar Medium
18.7 and 18.9 are good examples of this obscuration. Figure 18.14 shows another well-known, and particularly striking, example of such a cloud—the Horsehead Nebula in Orion. This curiously shaped finger of gas and dust projects out from the much larger dark cloud (called L1630) that fills the bottom half of the image and stands out clearly against the red glow of a background emission nebula. For refer-ence, the stars and bright emission nebulae lie in front of the dark cloud; the red glow that silhouettes the Horsehead lies behind and above it.
Absorption SpectraAstronomers first became aware of the true extent of dark interstellar clouds in the 1930s, as they studied the opti-cal spectra of distant stars. The gas in such a cloud absorbs some of the stellar radiation in a manner that depends on the cloud’s own temperature, density, and elemental abun-dance. The absorption lines thus produced contain infor-mation about dark interstellar matter, just as stellar absorp-tion lines reveal the properties of stars. (Sec. 4.1)
Because the interstellar absorption lines are produced by cold, low-density gas, astronomers can easily distinguish them from the much broader absorption lines formed in stars’ hot lower atmospheres. (Sec. 4.5) Figure 18.15(a) illustrates how light from a star may pass through several interstellar clouds on its way to Earth. These clouds need not be close to the star, and, indeed, they usually are not.
cloud clearly at radio wavelengths, providing an indispen-sible tool for the study of such objects. We will return to the subject of molecular emission from interstellar clouds in Section 18.5.
Figure 18.13 is a spectacular wide-field image of another dark dust cloud. Taking its name from a neigh-boring star system, Rho Ophiuchus, this dust cloud resides relatively nearby—about 170 pc from the Sun—making it one of the most intensely studied regions of star formation in the Milky Way. Pockets of heavy blackness mark regions where the dust and gas are especially concentrated and the light from the background stars is completely obscured. Measuring several parsecs across, the Ophiuchus cloud is only a tiny part of the grand mosaic shown in Figure 18.1. Note that this cloud, like most interstellar clouds, is very irregularly shaped. Note especially the long “streamers” of (relatively) dense dust and gas at upper left. By contrast, the bright patches within the dark regions are foreground objects—emission nebulae and groups of bright stars. Some of them are part of the cloud itself, where newly formed stars near the edge of the cloud have created “hot spots” in the cold, dark gas. Others have no connection to the cloud and just happen to lie along the line of sight.
Dark and dusty interstellar clouds are sprinkled throughout our Galaxy. We can study them at optical wave-lengths only if they happen to block the light emitted by more distant stars or nebulae. The dark outline of the L997 cloud in Figure 18.12(a) and the dust lanes visible in Figures
R I V U X G
Dust cloud
AntaresM4
Reflectionnebula
◀ FIGURE 18.13 Dark Dust Cloud The Ophiuchus dark dust cloud resides only 550 light-years away, surrounded by colorful stars and nebulae that are actually small illuminated parts of a much bigger, and invisible, molecular cloud engulfing much of the 6-degree-wide region shown. Many stages of star formation can be seen in this spectacular four-image mosaic. The dark cloud itself is “visible” only because it blocks light coming from stars behind it. Notice the cloud’s irregular shape, and especially its long “streamers” at upper left. The bright, giant star Antares, the (much more distant) star cluster M4, and a nearby blue reflection nebula are also noted. (R. Gendler/J. Misti/S. Mazlin)
The Ophiuchus dark dust cloud resides only 550 light-years away, surrounded by colorful stars and nebulae that are actually small illuminated parts of a much bigger, and invisible, molecular cloud engulfing much of the 6-degree-wide region shown.
The dark cloud itself is “visible” only because it blocks light coming from stars behind it.
The bright, giant star Antares, the (much more distant) star cluster M4, and a nearby blue reflection nebula are also noted.
10
‣Temperatura do MIS: desde alguns graus Kelvin até alguns 103 K (dependendo da proximidade de uma estrela ou fonte de radiação). Tem peratura média ~100 K; i.e. o meio interestelar é muito frio. ‣Densidade média ~ 10-24 g cm-3
Gás → 1 átomo/cm3 = 106 átomos/m3 Poeira → 1 partícula para 1012 átomos (1/1012)
‣Como então um material tão rarefeito e esparso pode bloquear a luz da estrela de modo tão eficaz? A resposta está na vastíssima extensão do meio interestelar.
11
Gás: ‣ O gás interestelar produz linhas de absorção quando a luz de uma estrela distante
interage com o gás ao longo da linha de visada. A absorção de radiação pelo gás depende de T, ρ e da abundância do gás.
‣ Como distinguir estas das linhas de absorção gás interestelar do da atmosfera da estrela?
-linhas da estrela são mais largas -linhas do gás interestelar são mais finas
‣ A abundância de alguns “metais” como C, O, Si, Mg, Fe é muito menor do que no sistema solar e nas estrelas. Razão possível: esses elementos são usados para formar poeira interestelar.
Poeira: → Observada no infravermelho. → Constituída de silicatos, C, Fe: sub-abundantes no gás! → Possivelmente também contém gelo “sujo” = agua congelada + amônia + metano + outras moléculas; semelhante a cauda de cometas
12
A luz pode ser absorvida ou espalhada pela poeira, desde que comprimento de onda da luz seja menor ou igual que a dimensão dos grãos:
λ ≤ dg A dimensão dos grãos é dg ~10-7 m, i.e. da mesma ordem que o comprimento de onda da luz visível. Portanto, regiões com poeira interestelar: ➡ são transparentes a ondas de radio e infravermelho: λ ≫ dg ➡ são opacas ao UV, raios X e raios gama: λ ≪ dg
Extinção interestelar
13
Nuvens densas de poeira são melhor observáveis no infravermelho ou em rádio (10-6m a 102 m). Para λs mais curtos (óptico, ultravioleta): elas são opacas, i.e. aparecem como nuvens escuras.Figura: Nuvem de poeira Barnard 68 no visível e no infravermelho.
SECTION 18.1 Interstellar Matter 451
Because the interstellar medium is more opaque to short-wavelength radiation than to radiation of longer wavelengths, light from distant stars is pref-erentially robbed of its higher frequency (“blue”) components. Hence, in addition to being generally diminished in brightness, stars also tend to appear redder than they really are. This effect, known as reddening, is conceptually similar to the process that produces spectacular red sunsets here on Earth.
(More Precisely 7-1)As illustrated in Figure 18.2(a), extinction and
reddening change a star’s apparent brightness and color. However, the patterns of absorption lines in the original stellar spectrum are still recognizable in the radiation reaching Earth, so the star’s spec-tral class can be determined. Astronomers can use this fact to study the interstellar medium. From a main-sequence star’s spectral and luminosity classes, astronomers learn the star’s true luminosity and color. (Secs. 17.5, 17.6) They then measure the degree to which the starlight has been affected by extinction and reddening en route to Earth, and this,
Scatteredlight
Dust cloud
TelescopeFrequency
Inte
nsity
Inte
nsity
Red lightslightlyreduced
Stellar absorptionlines still detectable
Blackbodyspectrum
Red BlueFrequency
Blue lightgreatly reduced
Stellarabsorption lines
(b)
(a)
R I V U X G
(c)
R I V U X G
StarNarrated FIGURE 18.2 Reddening (a) Starlight passing through a dusty region of space is both dimmed and reddened, but spectral lines are still recognizable in the light that reaches Earth. (b) This dusty interstellar cloud, called Barnard 68, is opaque
to visible light, except near its edges, where some light from background stars can be seen. The cloud spans about 0.5 light-year and lies about 520 light-years away. Frame (c) illustrates (in false color) how infrared radiation can penetrate Barnard 68. (ESO)
AN
IMAT
ION
/VID
EO I
nfra
red
View
of
Neb
ulae
AN
IMAT
ION
/VID
EO P
illar
s Be
hind
the
Dus
t
R I U X GV
in turn, allows them to estimate both the numbers and the sizes of interstellar dust particles along the line of sight to the star. By repeating these measurements for stars in many different directions and at many different distances from Earth, astronomers have built up a picture of the distribu-tion and overall properties of the interstellar medium in the solar neighborhood.
Reddening can be seen very clearly in Figure 18.2(b), which shows a type of compact, dusty interstellar cloud called a globule. (We will discuss such clouds in more detail in Section 18.3.) The center of this cloud, called Barnard 68, is opaque to all optical wavelengths, so starlight cannot pass through it. However, near the edges, where there is less intervening cloud matter, some light does make it through. Notice how stars seen through the cloud are both dimmed and reddened relative to those seen directly. Figure 18.2(c) shows the same cloud in the infrared part of the spectrum. Much more of the radiation gets through, but even here red-dening (or its infrared equivalent) can be seen.
14
SECTION 18.1 Interstellar Matter 451
Because the interstellar medium is more opaque to short-wavelength radiation than to radiation of longer wavelengths, light from distant stars is pref-erentially robbed of its higher frequency (“blue”) components. Hence, in addition to being generally diminished in brightness, stars also tend to appear redder than they really are. This effect, known as reddening, is conceptually similar to the process that produces spectacular red sunsets here on Earth.
(More Precisely 7-1)As illustrated in Figure 18.2(a), extinction and
reddening change a star’s apparent brightness and color. However, the patterns of absorption lines in the original stellar spectrum are still recognizable in the radiation reaching Earth, so the star’s spec-tral class can be determined. Astronomers can use this fact to study the interstellar medium. From a main-sequence star’s spectral and luminosity classes, astronomers learn the star’s true luminosity and color. (Secs. 17.5, 17.6) They then measure the degree to which the starlight has been affected by extinction and reddening en route to Earth, and this,
Scatteredlight
Dust cloud
TelescopeFrequency
Inte
nsity
Inte
nsity
Red lightslightlyreduced
Stellar absorptionlines still detectable
Blackbodyspectrum
Red BlueFrequency
Blue lightgreatly reduced
Stellarabsorption lines
(b)
(a)
R I V U X G
(c)
R I V U X G
StarNarrated FIGURE 18.2 Reddening (a) Starlight passing through a dusty region of space is both dimmed and reddened, but spectral lines are still recognizable in the light that reaches Earth. (b) This dusty interstellar cloud, called Barnard 68, is opaque
to visible light, except near its edges, where some light from background stars can be seen. The cloud spans about 0.5 light-year and lies about 520 light-years away. Frame (c) illustrates (in false color) how infrared radiation can penetrate Barnard 68. (ESO)
AN
IMAT
ION
/VID
EO I
nfra
red
View
of
Neb
ulae
AN
IMAT
ION
/VID
EO P
illar
s Be
hind
the
Dus
t
R I U X GV
in turn, allows them to estimate both the numbers and the sizes of interstellar dust particles along the line of sight to the star. By repeating these measurements for stars in many different directions and at many different distances from Earth, astronomers have built up a picture of the distribu-tion and overall properties of the interstellar medium in the solar neighborhood.
Reddening can be seen very clearly in Figure 18.2(b), which shows a type of compact, dusty interstellar cloud called a globule. (We will discuss such clouds in more detail in Section 18.3.) The center of this cloud, called Barnard 68, is opaque to all optical wavelengths, so starlight cannot pass through it. However, near the edges, where there is less intervening cloud matter, some light does make it through. Notice how stars seen through the cloud are both dimmed and reddened relative to those seen directly. Figure 18.2(c) shows the same cloud in the infrared part of the spectrum. Much more of the radiation gets through, but even here red-dening (or its infrared equivalent) can be seen.
‣ No caso da luz visível acontece o fenômeno de extinção interestelar (por absorção ou espalhamento) da luz.
‣ O fenômeno é tanto maior quanto menor o λ.
‣ A luz de estrelas mais distantes perde λs curtos (azuis). Logo, além de ter brilho menor, as estrelas parecem mais vermelhas.
‣ É o mesmo processo que produz o pôr do Sol avermelhado na Terra: a poeira do horizonte absorve os λs azuis e deixa passar os vermelhos.
‣ As linhas espectrais, no entanto, podem ser identificadas (ver figura).
15
‣ Devido à extinção interestelar, pode aumentar magnitude aparente (m). ‣ Portanto, é necessário incluir uma modificação na relação entre as magnitudes
aparente e absoluta de uma estrela: mV - MV = 5 log d - 5 + AV
onde AV é a absorção em magnitudes.
‣ Assim, se a extinção for alta, da ordem de alguns décimos de magnitude, o cálculo das distâncias será bastante afetado. Sem considerar a extinção interestelar, as distâncias das estrelas parecem maiores do que são na realidade.
‣ Além disso, a extinção afeta mais os comprimentos de onda mais curtos, produzindo um avermelhamento. Isso modifica a índice de cor (B-V).
16
Cor: B-V = mB-mV = -2,5 log (FB / FV)
FB > FV ⇨ B < V
[B-V] < 0
Estrela quente, azulada
FB < FV ⇨B > V
[B-V] > 0
Estrela fria, avermelhada
Fluxo
17
Correção da índice de cor‣ Linhas do espectro de absorção das estrelas: não são afetados pela poeira
presente no caminho entre a estrela e o observador
‣ Determinando tipo espectral e classe de luminosidade (pelo espectro) ⇨ T e
cor intrínseca da estrela (B-V)o.
‣ Verifica-se o quanto o valor observado [B-V] foi aumentado pela extinção. Definimos o avermelhamento como:
‣ E(B-V)=[B-V] – (B-V)o
‣ Correção do avermelhamento: parece razoável esperar que quanto mais extinção existe, o mais avermelhamento haverá. Na verdade, os estudos mostram que a extinção e vermelhidão são relacionadas por:
)VE(B3.1AV −=
18
O gás e a poeira interestelares podem se aglomerar em nuvens ou nebulosas. Dependendo das características, podemos ter:
• Nebulosa de emissão / reflexão • Nebulosa de Poeira• Nuvens moleculares
Nuvens interestelares
19
Nebulosas de reflexão: refletem a luz de estrela próxima em nossa direção.
Nebulosa de reflexão NGC1999, constelação de Orion.
20
Nebulosas de Emissão • As regiões HII, nuvens de gás
hidrogênio ionizado, ocorrem principalmente em volta de estrelas O e B pois estas emitem os fótons UV com energia acima de 13,6 eV. Quando os átomos de hidrogênio a absorvem, os elétrons ganham energia suficiente para se libertarem do núcleo, e o gás fica ionizado.
• Estas regiões portanto têm muitos íons de hidrogênio (prótons) e elétrons livres.
• Quando um próton captura um elétron livre, há emissão de radiação. As linhas do hidrogênio são emitidas quando o elétron passa, subsequentemente, pelos vários níveis de energia. Desta maneira, os fótons UV da estrela são degradados em fótons no visível pela região HII. Nebulosa de emissão de Orion, região HII
21
• A radiação emitida quando o elétron passa do nível n=3 para o n=2, em 658,3 nm, é dominante e causa a cor vermelha da região.
• Pelas medidas da largura das linhas no espectro da nebulosa obtém-se sua temperatura (nas proximidades da estrela): T ~ 8000K.
Nebulosa de emissão de Orion, região HII
22
Nuvens de poeira • Cerca de 90% do meio interestelar é composto por regiões escuras (sem estrelas ou
nebulosas brilhantes) chamadas de nuvens de poeira. • Elas são mais frias e mais densas que suas vizinhanças (103 a 106 vezes mais). • nmax = 1000 átomos/cm3 (densidade dos melhores vácuos em laboratório) • A maioria tem dimensões maiores que nosso sistema solar e apresentando formas
irregulares. Ocupam 2 - 3% do volume total do meio interestelar. • O principal componente dessas nuvens é o gás, como resto do meio interestelar.
Porém, há uma importante quantidade de poeira que absorve a luz visível. Por isso, as nuvens de poeira são melhor detectadas no infravermelho.
rho Oph no ótico rho Oph no IV
23
SECTION 18.3 Dark Dust Clouds 461
cases, the elemental abundances detected in interstellar clouds mirror those found in other astronomical objects—perhaps not surprising, since (as we will see in Chapter 19) interstellar clouds are the regions that spawn emission nebulae and stars.
PROCESS OF SCIENCE Check
4 How do astronomers use optical observations to probe the properties of dark dust clouds?
Each absorbs some of the stellar radiation in a manner that depends on its own temperature, density, velocity, and ele-mental abundance. Figure 18.15(b) depicts part of a typical spectrum produced in this way.
The narrow absorption lines contain information about dark interstellar clouds, just as stellar absorption lines reveal the properties of stars and nebular emission lines tell us about conditions in hot nebulae. By studying these lines, astrono-mers can probe the cold depths of interstellar space. In most
(b)(a)R I V U X G
1 light-year
▲ FIGURE 18.14 Horsehead Nebula (a) Located in the constellation Orion, not far from the Orion Nebula, this Horsehead Nebula is a striking example of a dark dust cloud, silhouetted against the bright background of an emission nebula. (b) A stunning image of the Horsehead, taken at highest resolution by the Very Large Telescope (VLT) in Chile. (Sec. 5.2) This nebular region is roughly 5000 light-years from Earth, in the constellation Orion. (Royal Observatory of Belgium; ESO)
(b)
Star
Stellar spectrum
Cloud 1
Cloud 2
Broad stellar absorption lines
Narrow absorptionlines from cloud 1
Fainter narrow absorptionlines from cloud 2
(a)
Inte
nsity
Frequency
Broad stellar lines
Narrow cloud lines
Real spectraare often messymixtures of spectrafrom many objectsalong the lineof sight.
◀ FIGURE 18.15 Absorption by Interstellar Clouds (a) Simplified diagram of some interstellar clouds between a hot star and Earth. Optical observations might show an absorption spectrum like that traced in (b). The wide, intense lines are formed in the star’s hot atmosphere; narrower, weaker lines arise from the cold interstellar clouds. The smaller the cloud, the weaker are the lines. The redshifts or blueshifts of the narrow absorption lines provide information on cloud velocities. The widths of all the spectral lines depicted here are greatly exaggerated for the sake of clarity.
AN
IMAT
ION
/VID
EO H
orse
head
Neb
ula
Horsehead Nebula (a) Located in the constellation Orion, not far from the Orion Nebula, this Horsehead Nebula is a striking example of a dark dust cloud, silhouetted against the bright background of an emission nebula. (b) A stunning image of the Horsehead, taken at highest resolution by the Very Large Telescope (VLT) in Chile. This nebular region is roughly 5000 light-years from Earth, in the constellation Orion.
24
‣O hidrogênio neutro (HI) emite uma linha espectral no comprimento de onda de 21 cm, que é usada para mapear a distribuição desse gás e que teve um papel chave na determinação da estrutura espiral da Galáxia. ‣ Os spins do elétron e do próton, no hidrogênio neutro em seu estado fundamental, podem ser paralelos (maior energia) ou opostos (menor energia). ‣A diferença de energia entre esses dois níveis é E = hν = 6×10−6 eV, que corresponde a uma frequência ν = E/h = 1420,4 MHz. ‣Portanto, a transição entre esses dois níveis de estrutura hiperfina dá origem a uma linha de comprimento de onda λ = c/ν = 21,049cm.
A linha de 21 cm
462 CHAPTER 18 The Interstellar Medium
low-energy state releases a photon with energy equal to the energy difference between the two levels.
Because that energy difference is very small, the energy of the emitted photon is very low. (Sec. 4.2) Consequently, the wavelength of the radiation is rather long—in fact, it is 21.1 cm, roughly the width of this book. That wavelength lies in the radio portion of the electromagnetic spectrum. Researchers refer to the spectral line that results from this hydrogen spin-flip process as 21-centimeter radia-tion. This spectral line provides a vital probe into any region of the universe containing atomic hydrogen gas. Figure 18.17 shows typical spectral profiles of 21-cm radio signals observed from several different regions of space. These tracings are the characteristic signatures of cold, atomic hydrogen in our Galaxy. Needing no visible star-light to help calibrate their signals, radio astronomers can observe any interstellar region that contains enough hydrogen gas to produce a detectable signal. Even the low- density regions between the dark clouds can be studied.
As can be seen in the figure, actual 21-cm lines are quite jagged and irregular, somewhat like nebular emission lines in appearance. The irregularities arise because there are usually numerous clumps of interstellar gas along any given line of sight, each with its own density, temperature, radial velocity, and internal motion. Thus, the intensity, width, and Doppler shift of the resultant 21-cm line vary from place to place. (Sec. 4.5) All these different lines are superimposed in the signal we eventually receive at Earth, and sophisticated computer analysis is generally required to disentangle them. The “average” figures quoted earlier for the temperatures (100 K) and densities (106 atoms/m3) of the regions between the dark dust clouds are based on 21-cm measurements. Observations of the dark clouds themselves using 21-cm radiation yield densities and temperatures in good agreement with those obtained by optical spectroscopy.
All interstellar atomic hydrogen emits 21-cm radia-tion. But if all atoms eventually fall into their lowest-energy
18.4 21-Centimeter RadiationA basic difficulty with the optical technique just described is that we can examine interstellar clouds only along the line of sight to a distant star. To form an absorption line, a background source must provide radiation to absorb. The need to see stars through clouds also restricts this approach to relatively local regions, within a few thousand parsecs of Earth. Beyond that distance, stars are completely obscured, and optical observa-tions are impossible. As we have seen, infrared observations provide a means of viewing the emission from some clouds, but they do not completely solve the problem because only the denser, dustier clouds emit enough infrared radiation for astronomers to study them in that part of the spectrum.
To probe interstellar space more thoroughly, we need a more general, more versatile observational method—one that does not rely on conveniently located stars and nebulae. In short, we need a way to detect cold, neutral interstellar matter anywhere in space through its own radiation. This may sound impossible, but such an observational technique does in fact exist. The method relies on low-energy radio emissions produced by the interstellar gas itself.
Electron SpinRecall that a hydrogen atom has one electron orbiting a single-proton nucleus. Besides its orbital motion around the central proton, the electron also has some rotational motion—that is, spin—about its own axis. The proton also spins. This model is analogous to a planetary system in which, in addition to the orbital motion of a planet about a central star, both the planet (electron) and the star (proton) rotate about their own axes. But bear in mind the crucial difference between plan-etary and atomic systems: A planet orbiting the Sun is free to move in any orbit and spin at any rate, but within an atom, all physical quantities, such as energy, momentum, and angular momentum (spin), are quantized—they are permitted to take on only specific, distinct values. (Sec. 4.2)
The laws of physics dictate that there are exactly two pos-sible spin configurations for a hydrogen atom in its ground state. The electron and proton can rotate in the same direc-tion, with their spin axes parallel, or they can rotate with their axes antiparallel (i.e., parallel, but oppositely oriented). Figure 18.16 shows these two configurations. The antiparallel configuration has slightly less energy than the parallel state.
Radio EmissionAll matter in the universe tends to achieve its lowest pos-sible energy state, and interstellar gas is no exception. A slightly excited hydrogen atom with the electron and pro-ton spinning in the same direction eventually drops down to the less energetic, opposite-spin state as the electron sud-denly and spontaneously reverses its spin. As with any other such change, the transition from a high-energy state to a
The emitted photoncarries away energyequal to the differencein the two spin states.
A hydrogen atom hasmore energy when itselectron and proton spinin the same direction.
Emittedphoton
ElectronProton
+
Antiparallel spins
Parallel spins
–
+ –
▲ FIGURE 18.16 Hydrogen 21-cm Emission A ground-level hydrogen atom changing from a higher-energy state (top) to a lower-energy state (bottom).
25
462 CHAPTER 18 The Interstellar Medium
low-energy state releases a photon with energy equal to the energy difference between the two levels.
Because that energy difference is very small, the energy of the emitted photon is very low. (Sec. 4.2) Consequently, the wavelength of the radiation is rather long—in fact, it is 21.1 cm, roughly the width of this book. That wavelength lies in the radio portion of the electromagnetic spectrum. Researchers refer to the spectral line that results from this hydrogen spin-flip process as 21-centimeter radia-tion. This spectral line provides a vital probe into any region of the universe containing atomic hydrogen gas. Figure 18.17 shows typical spectral profiles of 21-cm radio signals observed from several different regions of space. These tracings are the characteristic signatures of cold, atomic hydrogen in our Galaxy. Needing no visible star-light to help calibrate their signals, radio astronomers can observe any interstellar region that contains enough hydrogen gas to produce a detectable signal. Even the low- density regions between the dark clouds can be studied.
As can be seen in the figure, actual 21-cm lines are quite jagged and irregular, somewhat like nebular emission lines in appearance. The irregularities arise because there are usually numerous clumps of interstellar gas along any given line of sight, each with its own density, temperature, radial velocity, and internal motion. Thus, the intensity, width, and Doppler shift of the resultant 21-cm line vary from place to place. (Sec. 4.5) All these different lines are superimposed in the signal we eventually receive at Earth, and sophisticated computer analysis is generally required to disentangle them. The “average” figures quoted earlier for the temperatures (100 K) and densities (106 atoms/m3) of the regions between the dark dust clouds are based on 21-cm measurements. Observations of the dark clouds themselves using 21-cm radiation yield densities and temperatures in good agreement with those obtained by optical spectroscopy.
All interstellar atomic hydrogen emits 21-cm radia-tion. But if all atoms eventually fall into their lowest-energy
18.4 21-Centimeter RadiationA basic difficulty with the optical technique just described is that we can examine interstellar clouds only along the line of sight to a distant star. To form an absorption line, a background source must provide radiation to absorb. The need to see stars through clouds also restricts this approach to relatively local regions, within a few thousand parsecs of Earth. Beyond that distance, stars are completely obscured, and optical observa-tions are impossible. As we have seen, infrared observations provide a means of viewing the emission from some clouds, but they do not completely solve the problem because only the denser, dustier clouds emit enough infrared radiation for astronomers to study them in that part of the spectrum.
To probe interstellar space more thoroughly, we need a more general, more versatile observational method—one that does not rely on conveniently located stars and nebulae. In short, we need a way to detect cold, neutral interstellar matter anywhere in space through its own radiation. This may sound impossible, but such an observational technique does in fact exist. The method relies on low-energy radio emissions produced by the interstellar gas itself.
Electron SpinRecall that a hydrogen atom has one electron orbiting a single-proton nucleus. Besides its orbital motion around the central proton, the electron also has some rotational motion—that is, spin—about its own axis. The proton also spins. This model is analogous to a planetary system in which, in addition to the orbital motion of a planet about a central star, both the planet (electron) and the star (proton) rotate about their own axes. But bear in mind the crucial difference between plan-etary and atomic systems: A planet orbiting the Sun is free to move in any orbit and spin at any rate, but within an atom, all physical quantities, such as energy, momentum, and angular momentum (spin), are quantized—they are permitted to take on only specific, distinct values. (Sec. 4.2)
The laws of physics dictate that there are exactly two pos-sible spin configurations for a hydrogen atom in its ground state. The electron and proton can rotate in the same direc-tion, with their spin axes parallel, or they can rotate with their axes antiparallel (i.e., parallel, but oppositely oriented). Figure 18.16 shows these two configurations. The antiparallel configuration has slightly less energy than the parallel state.
Radio EmissionAll matter in the universe tends to achieve its lowest pos-sible energy state, and interstellar gas is no exception. A slightly excited hydrogen atom with the electron and pro-ton spinning in the same direction eventually drops down to the less energetic, opposite-spin state as the electron sud-denly and spontaneously reverses its spin. As with any other such change, the transition from a high-energy state to a
The emitted photoncarries away energyequal to the differencein the two spin states.
A hydrogen atom hasmore energy when itselectron and proton spinin the same direction.
Emittedphoton
ElectronProton
+
Antiparallel spins
Parallel spins
–
+ –
▲ FIGURE 18.16 Hydrogen 21-cm Emission A ground-level hydrogen atom changing from a higher-energy state (top) to a lower-energy state (bottom).
‣As colisões atômicas são suficientes para elevar a energia do HI para o estado excitado com spins paralelos. A energia necessária para produzir essa transição, pode ser obtida da energia térmica do gás kT; com T= 100 K . ‣O e l é t r o n p o d e d e c a i r espontaneamente para o estado de menor energia (spins anti-paralelos) emitindo um fóton na frequência de rádio.
26
‣ FIGURA: Linhas de 21 cm típicas observadas em várias regiões diferentes do meio interestelar.
‣ Por causa da al ta abundância de hidrogênio, apesar da longa vida média, ela é observada em todas as direções do céu.
‣ O s p i c o s n ã o o c o r r e m t o d o s n o comprimento de onda de exactamente 21,1 cm (1420 MHz), porque o gás em nossa Galáxia está em movimento em relação à terra.
‣ Essa emissão em rádio atravessa facilmente o meio interestelar, mesmo regiões com muita poeira e chega até nós.
SECTION 18.5 Interstellar Molecules 463
CONCEPT Check
4 Why is 21-cm radiation so useful as a probe of galactic structure?
18.5 Interstellar MoleculesIn some particularly cold (typically, 10–20 K) interstellar regions, densities can reach as high as 1012 particles/m3. Until the late 1970s, astronomers regarded these regions simply as abnormally dense interstellar clouds, but it is now recognized that they belong to an entirely new class of interstellar matter. The gas particles in these regions are not in atomic form at all; they are molecules. Because of the predominance of mol-ecules in these dense interstellar regions, they are known as molecular clouds. They literally dwarf even the largest emis-sion nebulae, which were previously thought to be the most massive residents of interstellar space.
Molecular Spectral LinesAs noted in Chapter 4, much like atoms, molecules can become excited through collisions or by absorbing radiation.
(Sec. 4.4) Furthermore, again like atoms, molecules even-tually return to their ground states, emitting radiation in the process. The energy states of molecules are much more com-plex than those of atoms, however. Once more like atoms, mol-ecules can undergo internal electron transitions, but unlike atoms, they can also rotate and vibrate. They do so in specific ways, obeying the laws of quantum physics. Figure 18.18 depicts a simple molecule rotating rapidly—that is, a molecule in an excited rotational state. After a length of time that depends on its internal makeup, the molecule relaxes back to a slower rotational rate (a state of lower energy). This change causes a photon to be emitted, carrying an energy equal to the energy difference between the two rotational states involved. The energy differences between these states are generally very small, so the emitted radiation is usually in the radio range.
configuration, then why isn’t all the hydrogen in the Galaxy in the lower energy state by now? Why do we see 21-cm radia-tion today? The answer is that the energy difference between the two states is comparable to the energy of a typical atom at a temperature of 100 K or so. As a result, atomic collisions in the interstellar medium are energetic enough to boost the electron into the higher energy configuration and so main-tain comparable numbers of hydrogen atoms in either state. At any instant, any sample of interstellar hydrogen will con-tain many atoms in the upper level, so conditions will always be favorable for 21-cm radiation to be emitted.
Of great importance, the wavelength of this character-istic radiation is much larger than the typical size of inter-stellar dust particles. Accordingly, 21-cm radiation reaches Earth completely unscattered by interstellar debris. The opportunity to observe interstellar space well beyond a few thousand parsecs, and in directions lacking background stars, makes 21-cm observations among the most impor-tant and useful in all astronomy. We will see in Chapters 23 through 25 how such observations are indispensable in allowing astronomers to map out the large-scale structure of our Galaxy and many others.
Inte
nsity
Frequency1420 MHz(Wavelength = 21.1 cm)
▲ FIGURE 18.17 21-cm Lines Typical 21-cm radio spectral lines observed from several different regions of interstellar space. The peaks do not all occur at a wavelength of exactly 21.1 cm, corresponding to a frequency of 1420 MHz, because the gas in our Galaxy is moving with respect to Earth.
This is a spinning formaldehydemolecule, H2CO.
H
H
H
O
O Emittedphoton
CC
H
▲ FIGURE 18.18 Molecular Emission As a molecule changes from a rapid rotation (left) to a slower rotation (right), a photon is emitted that can be detected with a radio telescope. The lengths of the curved arrows are proportional to the spin rate of the molecule.
27
Observação de nuvens interestelares
‣ Regiões HII: gás ionizado por estrelas quentes; as temperaturas chegam até 104 K (nebulosas de emissão). Podem ser detectadas no óptico.
‣ Regiões HI: gás neutro, T = 50 a 100 K, são detectadas pela linha de 21 cm.
28
‣ As nuvens moleculares estão compostas predominantemente de moléculas. ‣ As primeiras moléculas interestelares foram descobertas em 1937-1938, na
forma de metilidina CH, CH+, e cianogênio CN, aparentes nos espectros de algumas estrelas, mas causadas por absorção interestelar.
‣ Hidrogênio molecular H2 foi descoberto no início dos anos 1970, junto com monóxido de carbono CO. Muitos outros tipos de moléculas têm sido encontradas desde então, desde amônia NH3, até as mais complexas como etanol C2H5OH.
‣ Baseado principalmente nas observações de CO, nota-se que as moléculas estão concentradas em nuvens moleculares, com massas de poucas vezes até um milhão de massas solares, e se estendem de alguns até cerca de 600 anos-luz. As temperaturas são da ordem de 10 K.
‣ As estrelas se formam nas partes mais densas destas nuvens moleculares.
Nuvens moleculares
29
‣Moléculas e poeira sempre aparecem associadas: Por que?
‣A presença de poeira nas nuvens tem um efeito de preservar as moléculas contra radiação de altas energias, que poderia destruí-las.
‣Poeira também pode agir como catalisadora na formação de moléculas, já que propiciam um local para átomos juntarem-se e dissipar a energia da reação que de outra forma poderia dissociar moléculas recém formadas.
30
Observação das nuvens moleculares ‣O principal conteúdo destas nuvens é H2 ; mas não emite nem absorve radiação rádio. Embora o hidrogênio molecular produza linhas no ultravioleta, o gás e poeira existentes nas nuvens moleculares torna a extinção ultravioleta muito grande, dificultando a medida do H2. ‣ Por outro lado, a linha de 21cm só é emitida pelo H atômico e não molecular. ‣Portanto, é necessário observar outras moléculas (ex.: CO, HCN, NH3, H2O, H2CO): mas são milhões a bilhões de vezes menos abundantes que o H2. A emissão ocorre quando há mudanças no estado rotacional das moléculas. As diferenças de energia são pequenas, por isso a emissão ocorre principalmente em ondas rádio.
31
• Nuvens moleculares não se encontram isoladas.
• Formam complexos com dimensões de até ~100 pc.
• Os maiores complexos contêm gás suficiente para formar milhões de estrelas como o Sol.
• Cerca de 1000 complexos desse tipo são encontrados em nossa Galáxia.
32
Formação de estrelas
•Vimos em um capítulo anterior que estrelas formam-se a partir do colapso gravitacional de nuvens (escuras ou moleculares). •A nuvem deve ser massiva e fria o bastante para que FG > FP. •A poeira constitui uma fração desprezível da massa total da nuvem. No entanto, o conteúdo de poeira é importante para resfriar o gas à medida que contrai e mais tarde para formação de planetas. •Cloud Fragmentation: As an interstellar cloud contracts, gravitational instabilities cause it to fragment into smaller pieces. The pieces themselves continue to fall inward and fragment, eventually forming many tens or hundreds of individual stars.
SECTION 19.2 The Formation of Stars Like the Sun 473
when a nearby O- or B-type star forms and ionizes its sur-roundings, squeezes a cloud beyond the point where pres-sure can resist gravity’s inward pull. (Secs. 16.2, 17.5) Or perhaps the cloud’s supporting magnetic field leaks away as charged particles slowly drift across the confining field lines, leaving the gas unable to support its own weight (More Pre-cisely 19-1).
Whatever the cause, theory suggests that once the collapse begins, fragmentation into smaller and smaller clumps of matter naturally follows, as gravitational instabilities continue to operate in the gas. As illustrated in Figure 19.3, a typical cloud can break up into tens, hundreds, or even thousands, of fragments, each imi-tating the shrinking behavior of the parent cloud and contracting ever faster. The whole process, from a single stable cloud to many collapsing fragments, takes a few million years.
In this way, depending on the precise conditions under which fragmentation takes place, an interstellar cloud can produce either a few dozen stars, each much larger than
Note again the time scales involved in these stages—even the shortest spans a thousand human generations. Astronomers have not gained this insight by watching a single cloud or group of clouds evolve from start to finish. Rather, they combine theory and observation to refine a still evolving mathematical model of how stars form.
Stage 1: An Interstellar CloudThe first stage in the star-formation process is a dense interstellar cloud—the core of a dark dust cloud or perhaps a molecular cloud. These clouds are truly vast, sometimes spanning tens of parsecs (1014–1015 km) across. Typical temperatures are about 10 K throughout, with a density of perhaps 109 particles/m3. Stage-1 clouds contain thou-sands of times the mass of the Sun, mainly in the form of cold atomic and molecular gas. (The dust in a stage-1 cloud both cools the cloud as it contracts and plays a cru-cial role in planet formation, but it constitutes a negligible fraction of the total mass of the cloud.) (Sec. 15.2)
Despite their low internal tem-peratures, most observed dark inter-stellar clouds seem to have enough internal pressure to support them-selves against the force of gravity.
(More Precisely 8-1) However, if such a cloud is to be the birthplace of stars, it must become unstable, start to collapse under its own gravity, and eventually break up into smaller pieces. Most astronomers think that the process of star formation is trig-gered when some external event, such as the shock of a nearby stellar explo-sion or the pressure wave produced
In reality, no interstellar cloud ever fragments this neatly; it’s usually a mess.
▲ FIGURE 19.3 Cloud Fragmentation As an interstellar cloud contracts, gravitational instabilities cause it to fragment into smaller pieces. The pieces themselves continue to fall inward and fragment, eventually forming many tens or hundreds of individual stars.
Stage Approximate Time to Next Stage
(yr)
Central Temperature
(K)
Surface Temperature
(K)
Central Density
(particles/m3)
Diameter* (km)
Object
1 2 * 106 10 10 109 1014 Interstellar cloud
2 3 * 104 100 10 1012 1012 Cloud fragmentCloud fragment/protostar
3 105 10,000 100 1018 1010
4 106 1,000,000 3000 1024 108 Protostar
5 107 5,000,000 4000 1028 107 Protostar
6 3 * 107 10,000,000 4500 1031 2 * 106 Star
7 1010 15,000,000 6000 1032 1.5 * 106 Main-sequence star
* Round numbers; for comparison, recall that the diameter of the Sun is 1.4 * 106 km, whereas that of the solar system is roughly 1.5 * 1010 km.
TABLE 19.1 Prestellar Evolution of a Solar-Type Star
33
Nebulosa Trífide M20, na constelação de Sagitário, a 9000 anos-luz de distância, exemplo de um berçário de estrelas .
34
Região de M20 (Trifide), mostra três estágios da formação estelar:
•Nuvem progenitora
•Fragmentos colapsando
•Nebulosas de emissão contendo estrelas recém formadas
480 CHAPTER 19 Star Formation
obscured regions labeled A and B, where the molecu-lar emission of radio energy is strongest, are such denser, warmer fragments. Here, the total gas density is observed to be at least 109 particles/m3, and the temperature is about 100 K. The Doppler shifts of the radio lines observed in the vicinity of region B imply that this portion of M20, labeled “contracting fragment” in the figure, is infalling. Recent infrared observations (Figure 19.8c) reveal the candidate protostars themselves, identified by the warmth of their growing embryos tucked inside. Less than a light-year across, the region has a total mass over a thousand times the mass of the Sun—considerably more than the mass of M20 itself. The region lies somewhere between stages 1 and 2 of Table 19.1.
The third star-formation phase shown in Figure 19.8 is M20 itself. The glowing region of ionized gas results directly from a massive O-type star that formed there within the past million years or so. Because the central star is already fully formed, this final phase corresponds to stage 6 or 7 of our evolutionary scenario.
Evidence of Cloud FragmentsOther parts of our Milky Way Galaxy provide sketchy evi-dence for prestellar objects in stages 3 through 5. The Orion complex, shown in Figure 19.9, is one such region about 1400 light-years away. Lit from within by several O-type stars, the bright Orion Nebula is partly surrounded by a vast molecular cloud that extends even beyond the roughly 10 * 30-light-year region bounded by the photograph in Figure 19.9(b).
The Orion molecular cloud harbors several smaller sites of intense radiation emitted by molecules deep within
The interstellar clouds in and around M20 thus provide tentative evidence of three distinct phases of star formation, as shown in Figure 19.8. The huge, dark molecular cloud surrounding the visible nebula is the stage-1 cloud. Both its density and its temperature are low—about 108 particles/m3 and 20 K, respectively. Greater densities and temperatures typify smaller regions within this large cloud. The totally
R I V U X G
R I V U X G R I V U X G
Densefragments
Excitingstar
B
A
Parentcloud
Contractingfragment
(a)
(b)
5 pc
(c)
Nebula
◀ FIGURE 19.8 Star Formation Phases (a) The M20 region shows observational evidence for three broad phases in the birth of a star. The parent cloud is stage 1 of Table 19.1. The region labeled “contracting fragment” likely lies between stages 1 and 2. Finally, the emission nebula (M20 itself) results from the formation of one or more massive stars (stages 6 and 7). (b) A close-up (including Hubble inlays) of the area near region B outlines (in drawn ovals) especially dense knots of dusty matter. (c) A Spitzer Telescope infrared image of the same scene reveals those cores thought to be stellar embryos (arrows). (AURA; NASA)
35
Pilares Gasosos (Nebulosa da Águia)
36
37
Nebulosa Tarântula (30 Doradus) rico berçário estelar:
A. Neb. do Cone, em Monoceros.
B. Antares: luz espalhada por grãos de diferentes tamanhos
C. CG4 neb. de reflexão
D. IC2944 neb. de emissão + Glóbulos de Bok
