1 / 25

Chapter 3

Chapter 3. Stars: Radiation.  Nick Devereux 2006. Revised 2007. Blackbody Radiation.  Nick Devereux 2006. The Sun (and other Stars) radiate like Blackbodies.  Nick Devereux 2006. The Planck Function. I  = 2 h  3 W m -2 Hz -1 sterad -1 c 2 (e h  / kT - 1). Where.

geoffrey-hogan
Download Presentation

Chapter 3

An Image/Link below is provided (as is) to download presentation Download Policy: Content on the Website is provided to you AS IS for your information and personal use and may not be sold / licensed / shared on other websites without getting consent from its author. Content is provided to you AS IS for your information and personal use only. Download presentation by click this link. While downloading, if for some reason you are not able to download a presentation, the publisher may have deleted the file from their server. During download, if you can't get a presentation, the file might be deleted by the publisher.

E N D

Presentation Transcript

  1. Chapter 3 Stars: Radiation  Nick Devereux 2006 Revised 2007

  2. Blackbody Radiation  Nick Devereux 2006

  3. The Sun (and other Stars) radiate like Blackbodies  Nick Devereux 2006

  4. The Planck Function I = 2 h 3W m-2 Hz-1 sterad-1 c2 (eh/kT - 1) Where h is Planck’s constant 6.626 x 10-34 J s k is Boltzmanns constant 1.380 x 10-23 J/K c is the speed of light 2.998 x 108 m/s T is the temperature in K  is the frequency in Hz and Iis the Specific Intensity  Nick Devereux 2006

  5. Ivs. I It is important to know which type of plot you are looking at Ior I.  Nick Devereux 2006

  6. Transferring from Ito I • Id= Id(equivalent energy) • Since c =  •  = c/  Thus, d = -c d 2  Nick Devereux 2006

  7. Then, I = Id d I = - 2 h 3c c2 (eh/kT - 1) 2 I = - 2 h c2 W m-2 m-1 sterad-1 5 (eh /kT - 1)  Nick Devereux 2006

  8. Wiens Law • Differentiating Ileads to Wien’s Law, • max T = 2.898 x 10-3 • Which yields the peak wavelength, max (m). • for a blackbody of temperature, T.  Nick Devereux 2006

  9. Blackbody Facts Blackbody curves never cross, so there is no degeneracy. The ratio of intensities at any pair of wavelengths uniquely determines the Blackbody temperature, T. Since stars radiate approximately as blackbodies, their brightness depends not only on their distance, but also their temperature and the wavelength you observe them at.  Nick Devereux 2006

  10. Temperature Determination To measure the temperature of a star, we measure it’s brightness through two filters. The ratio of the brightness at the two different wavelengths determines the temperature. The measurement is independent of how far away the star is because distance reduces the brightness at all wavelengths by the same amount.  Nick Devereux 2006

  11. Filters and U,B,V Photometry Filters transmit light over a narrow range of wavelengths  Nick Devereux 2006

  12. The Color of a Star is Related to it’s Temperature  Nick Devereux 2006

  13. Color Index A quantitative measure of the color of a star is provided by it’s color index, defined as the difference of magnitudes at two different wavelengths. mB – mV = 2.5 log {fV/fB} + c The constant sets the zero point of the system, defined by the star Vega which is a zero magnitude star. Magnitudes for all other stars are measured with respect to Vega.  Nick Devereux 2006

  14. Dealing with the constant In the basic magnitude equation, there is a constant, c, which I can now tell you is equivalent to mo = -2.5 log (the flux of the zero magnitude star Vega). So, for a star of magnitude m* we can write m* - mo = 2.5 log {fo/f*} Note: There is no constant ! In this equation mo = 0 of course because it is the magnitude of a zero magnitude star. However, the flux of the zero magnitude star, fo is not zero, as you can see on the next slide.  Nick Devereux 2006

  15. Zero Magnitude Fluxes Filter  (m) F (W/cm2 m)F (W/m2 Hz) U 0.36 4.35 x 10-12 1.88 x 10-23 B 0.44 7.20 x 10-12 4.44 x 10-23 V 0.55 3.92 x 10-12 3.81 x 10-23 1 Jansky (Jy) = 1 x 10-26 W/m2 Hz  Nick Devereux 2006

  16. Calculating Fluxes Now you know what the fluxes are for a zero magnitude star, fo, you can convert the magnitudes for any object in the sky (stars, galaxies, etc) into real fluxes with units of Wm-2 Hz-1, at any wavelength using this equation! m* = 2.5 log {fo/f*}  Nick Devereux 2006

  17. Vega ( also known as -Lyr) Vega has a temperature ~ 10,000 K, so it is a hot star. Vega is the zero magnitude star, it’s magnitude is defined to be zero at all wavelengths. Be aware - This does not mean that the flux is zero at all wavelengths!! Magnitudes for all other stars are measured with respect to Vega, so stars cooler than Vega have B-V > 0, and stars warmer than Vega have B-V < 0.  Nick Devereux 2006

  18. Color and Temperature The B-V color is directly related to the temperature.  Nick Devereux 2006

  19. Bolometric Magnitudes ( MBol ) When we measure Mv for a star, we are measuring only the small part of it’s total radiation transmitted in the V filter. To get the Bolometric magnitude, MBol which is a measure of the stars total output over all wavelengths, we make use of a Bolometric Correction (BC). So that, MBol = Mv + BC The BC depends on the temperature of the star because Mv includesdifferent fractions of MBol depending on the temperature (see Appendix E). Question: The BC is a minimum for 6700K – Why?  Nick Devereux 2006

  20. The Sun The Sun has a BC = -0.07 mag and a bolometric magnitude, Mbol(sun) = +4.75 mag, and an effective temperature = 5800K.  Nick Devereux 2006

  21. Spectral Types There is a system for classifying stars that involves letters of the alphabet; O,B,A,F,G,K,M. These letters order stars by Temperature, with O being the hottest, and M the coldest. Our Sun is a G type star. Vega is an A type star. The letter sequence is subdivided by numbers 0 to 5, with 0 being the hottest. So a BO star is hotter than a B5 star.  Nick Devereux 2006

  22. Luminosity Classes Stars are also subdivided on the basis of their evolutionary status, identified by the Roman numerals I,II,III,IV and V. There will be more about this later. Stars spend most of their lives on the main sequence, luminosity class V. The Sun is a GOV.  Nick Devereux 2006

  23. Stellar Luminosity The Stellar Luminosity is obtained by integrating the Planck function over all wavelengths, and eliminating the remaining units (m-2 sterad–1), by multiplying by 4π D2, the spherical volume over which the star radiates, and the , the solid angle the star subtends, to obtain L = 4π R2 T4 W Where R is the radius of the star, T is the stellar temperature, and is the Stefan-Boltzmann constant = 5.67 x 10-8 W m-2 K-4  Nick Devereux 2006

  24. Relating Bolometric Magnitude to Luminosity The bolometric magnitudes for any object, Mbol* , may be compared with that measured for the Sun, Mbol, to determine the luminosity of the object, L* in terms of the luminosity of the Sun, L○. Mbol - Mbol* = 2.5 log{ L* / L }  Nick Devereux 2006

  25. You now know how to measure the luminosity and temperature of stars. Next, we need to find their masses. Once we have done that we can plot a graph like the one on the left. Stars populate a narrow range in this diagram with the more massive ones having higher T and L. Understanding the reason for this trend will lead us to an understanding of the physical nature of stars. Where we are going …..  Nick Devereux 2006

More Related