Ig2 ANNUAL KEPOllT SMITUSONIAN INSTITUTION, 1937 



telescope. From each line of the flash spectrum information is 

 obtained in three separate and distinct directions: 1, wave lengths; 

 2, intensities; and 3, heights. 



With the brief exposures available at eclipse time it is only possible 

 to use a dispersion much smaller than that employed in daily solar 

 work, with the result that the accuracy of wave-length determinations 

 is not sufficiently high to ascertain systematic dift'erences between 

 eclipse wave lengths and those taken under ordinary solar conditions. 

 Accordingly, chromospheric wave lengths can serve no other purpose 

 than the identification of the lines from comparisons with Rowland's 

 Tables in order to determine the gaseous element whence the sj)ectral 

 lines originate. 



The most characteristic difference between the chromospheric and 

 the ordinary solar spectrum is found in the relative intensities of the 

 lines. In the eclipse spectra the helium lines, including D3 discovered 

 in 1SG8, are of great strength, whereas in the Fraunhofer spectrum 

 they are entirely missing. In the eclipse spectra the hydrogen lines 

 are of great prominence, and 32 lines of the Balmer series have been 

 observed, whereas in the Fraunhofer spectrum only 4 hydrogen lines 

 are observed. Compared side by side, the spectra on plate 7 seem 

 to belong to stars of two different types rather than to the same object 

 under different conditions. 



A glance at the chromospheric spectra taken without slit shows that 

 the strongest lines are of the greatest lengths on the photograph and 

 hence the vapors extend to the greatest heights above the sun's sur- 

 face. The heights from eclipse spectra in the capable hands of Saha 

 provided him with indispensable information from which he was able 

 to derive temperature and pressure conditions in the solar atmosphere. 

 The theory of ionization developed a decade and a half ago as a result 

 of heights from eclipse spectra has fairly revolutionized the science 

 of astrophysics. Saha's theory has furnished an explanation of the 

 causes of the differences in intensities of lines in spectra of stars of 

 different types, has given a means of deriving the temperatures of 

 the stars and has given measures of the distances of the stars by 

 their spectroscopic parallaxes. 



Starting out with Saha's theory, new information gleaned about the 

 structure of the atom has revolutionized the method of attacking 

 solar problems. Practically all the prominent lines in the spectrum 

 of the sun have been assigned to multiplet groups with known excita- 

 tion potentials measured in electron-volts, the arbitrary intensity 

 scale of Rowland has been submitted to calibration tests which have 

 revealed that the intensities depend on the number of atoms engaged 

 in the formation of the spectral lines. From the weakest Fe lines 

 perceptible in the solar spectra to the strongest, the number of atoms 

 involved increases about 1,000,000 times. 



