The Hertzsprung–Russell diagram

6.9 The Hertzsprung–Russell diagram

In the early 1900s, Ejnar Hertzsprung in Denmark and Henry Russell at Princeton University in the USA independently plotted a graph of stellar luminosity against temperature which has become known as the Hertzsprung–Russell diagram (H–R diagram) (Figure 6.6). Both axes are logarithmic. The x axis represents tempera- ture. (Note: it increases to the left of the plot.) Instead of using temperature to defi ne the ordinate, both spectral type and the (B-V) colour index can be used as

Figure 6.6 The Hertzsprung–Russell Diagram.

Introduction to Astronomy and Cosmology

both are directly related to temperature. Henry Russell’s fi rst diagram, published in the journal Nature in 1914, used the spectral type. An observer might use the B-V index and a theorist would be likely to use temperature, regarding this as being the most fundamental. On the plot shown in Figure 6.6, both the spectral type and the temperature are given.

The vertical axis is a measure of brightness. Here, a theoretician might well use

a star’s luminosity compared with the Sun, whilst an observer would use abso- lute magnitude. Russell used the absolute magnitude for this ordinate, but this author prefers the use of luminosity as giving a better understanding of the rela- tive brightness of the stars. When plots are used using the colour index and abso- lute magnitude, they are often called a colour–magnitude diagram rather than an H–R diagram (Figure 6.7).

To provide a totally true diagram, all stars within a volume of space suffi ciently large to include a reasonable number of the very rare 0 type stars would need to be plotted – this would be known as a complete sample. This is very diffi cult to achieve as, in this case, the volume is so large that the majority of the faintest stars, those of Type M, would be too faint to be detected, so the plot presented is not a com- pletely true refl ection of the relative numbers of the different types of stars. It does, however, show where a representative sample of stars lie on the H–R diagram.

6.9.1 The main sequence

The main sequence is an S-shaped region that extends from the top left (very bright, high surface temperature O type stars) down to the bottom right (faint, low surface temperature M type stars) of Figure 6.6. It is the region which includes between 80% and 90% of all stars. The stars at the lower right of the main sequence are called red dwarfs as their luminosity is far less than that of our Sun.

6.9.2 The giant region

Above the main sequence on the right-hand side of Figure 6.6 is an area of bright stars whose colours range through yellow, orange and red. As these stars are very bright they are called giant stars , such as the star Aldebaran, in Taurus, which is called a red giant (though they actually appear orange in colour). At the top of Figure 6.6 is a region, extending from the blue through to the red, of exceedingly bright stars which are called supergiants . Betelgeuse, in Orion, at the extreme top right of the plot is a red supergiant. In contrast, at the extreme top left of the plot is the star Rigel, the brightest star in Orion, whose luminosity we calculated to be ∼45 000 times that of the Sun. This is termed a blue supergiant.

The Properties of Stars

Figure 6.7 Colour–magnitude diagram of the 16 631 brightest stars from the Hipparcos catalogue showing the number of stars in given colour–magnitude cells. (Note that this greatly overemphasizes the number of bright stars as they can be seen at great distances.) Image: Hipparcos Space Astrometry Mission, ESA.

Introduction to Astronomy and Cosmology

6.9.3 The white dwarf region

Below the main sequence lies a region in which white dwarf stars are found. (They encompass a wide surface temperature range and are not necessarily white.) The companion to Sirius, in Canis Major, is a white dwarf. As we will see in Chapter 7 white dwarfs are the remnants of stars like our Sun. They are very small, about the size of the Earth, so even those with very high surface temperatures are not very luminous.

6.9.4 Pressure broadening

As a red dwarf and a red giant or red supergiant can all have the same surface temperature, one might well ask how they can be distinguished. Obviously their measured luminosity is a major factor but, in addition, their spectra appear differ- ent. It is possible to measure the width of spectral lines and it is found that those of the giant stars are less than those of the dwarf stars. As we will shortly see, the outer envelopes of red giant stars are very tenuous with a very low gas pressure whereas the pressure in the atmosphere of M type dwarfs is far higher. The width of an emitted spectral line in gas of higher pressure is broadened in comparison with that in a low pressure due to a phenomenon called pressure broadening : an atom which is able to emit a photon of energy unhindered will produce a narrow band of emitted frequencies but if this atom collides with another atom during the emission process (which will happen far more often at higher pressures), the emitted wave train is shortened in comparison, and this gives rise to a wider range of frequencies – a broader spectral line. This same effect is seen in the absorption lines of stars with red giants having narrower spectral lines than red dwarfs.

The reason why most stars are seen to lie on the main sequence is simply because this is where stars spend the majority of their life as stable objects produc- ing energy by the nuclear fusion of hydrogen to helium. As stars evolve, their posi- tion in the H–R diagram changes and a star is said to move along an evolutionary track across the H–R diagram. The tracks of stars of different stellar types will be discussed in Chapter 7. The main sequence is not a line but, as shown, is a band across the H–R diagram. This is because, as they age, stars become somewhat more luminous with an increased surface temperature. They thus move some- what up and to the left of the diagram during their hydrogen burning phase. The giant phase in the life of a star is relatively brief which is why we see far fewer stars of this type. The white dwarfs are the fi nal states of many stars and gradually cool over billions of years thus moving down and to the right of the H–R diagram. Over time, as more of the stars in our Galaxy come to the end of their lives, their numbers will increase relative to those on the main sequence.

The Properties of Stars