The Light Curve of XX Cygni As a Contribution to the Study of Cepheid Variation (1915)(en)(5s)


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Gu, J. Hong, I. McMichael, R. Saxena, and F. Sincerbox, P. Wimmer, and G.

Period Changes in SX Phoenicis Stars. III. XX Cygni

Coufal, D. Sincerbox, eds. Berlin, Germany: Springer-Verlag, , pp. F9 Giant G5 9. Aur Supergiant K5 Main Seq. B9 F9 Giant G9 AO Giant G8 2.

UZ Cyg Main Seq. A1 Giant K1 A3 Dwarf G65 6. Successive columns contain the name of the star, the maximum brightness, the amplitudes, the period, and the two spectra. The last col- umn shows the difference in spectra, in units of the spectral class. Ai Aa Period Spectra Eel. AS SY And m Much uncertainty in classifying the fainter spectrum.

The secondary is a duplicate of the primary except in in- tensity. The spectra are identical and nearly of the same in- tensity. Both stars show the helium lines, but with unequal intensity. Spectra well separated. The lines of the secondary are about 0. The lines are of good quality. The separation has been made only on one plate; spectro- scopically this system has been poorly investigated. The uncertainty of the second spectrum is due to the pres- ence of the primary spectrum on the same plate. The chief features of the spectrum, , and , are fairly narrow emission lines, characteristic of a small group of absorption O type stars.

The second spectrum is very faint, appearing to be of the same type as that of the primary star. None of the three plates at primary minimum is entirely sat- isfactory, as they were taken under poor conditions, but they show that the fainter star is of type KO and probably belongs to the dwarf branch. Enhanced line spectrum similar to that of a Cygni. HP and Hv have bright edges. The faint lines of the secondary spectrum, approximately of the same type as the primary, have been seen and measured on several plates.

Both appear to be of the same spectrum. The spectrum resembles closely that of AO Cas, but is of slightly earlier type. The strongest lines are the Balmer series. Neutral helium , is stronger than ionized helium , Doubly ionized carbon , trebly ionized silicon , and singly ionized nitrogen are the next strongest lines, followed by the weaker lines of neutral hel- ium. The second spectrum, which appears to be of the same spectral type, is exceedingly faint. The spectrum of the fainter is plainly visible.

The spectra of both components are visible. The lines are double, of approximately equal strength, and very diffuse and shallow. One observation has been secured during the secondary mini- mum and the spectrum is typical B9n. Hp is a weak emission line, and the spectrum is continuous at Hy. The spectra of components are unusual.

The absorption spec- trum pertaining to the more massive component is somewhat stronger than that of the less massive component, approxi- mately in the ratio of The stars belong to the small group of M type dwarfs which show H and K and the Balmer series as emission lines.

All the bright lines give the same velocities as the absorption lines of the corresponding spec- trum.

Period Changes in SX Phoenicis Stars. III. XX Cygni - IOPscience

When the velocities are such as to separate the lines, the red member of the pair emission lines is, with one or two exceptions, the stronger. The lines of the second spectrum are only about one-half as strong as those of primary, and many are poorly defined or not seen at all on some plates. The spectrum of the star eclipsed at primary minimum is stronger. One plate taken at the primary mini- mum. The spectra resemble closely those of a Virginis. The spec- trum of the secondary component is exactly similar to that of the primary, but the lines are of the vanishing order of intensity, and exceedingly difficult to measure.

The spectrum of the secondary at the primary minimum contains bright hydrogen lines. The lines of the brighter component are about three times as strong as those of the primary, which, though the more mass- ive, is in this system the fainter component. The lines of both components are very diffuse. The Class K spectrum has sharp emission lines, H and K of Call, superposed on broad absorption lines.

On some spectra the lines appear to be double, as though two stars of about the same spectral class were present. Second spectrum is extremely difficult to measure. In the secondary spectrum at primary minimum H 3 is bright and fairly strong. The spectrum is continuous at Hy. The lines of the primary are slightly stronger than those of the secondary.

The spectrum is characterized by extremely wide, shallow lines. The spectrum at primary minimum is A3. Second spectrum is practically B3. The spectrum of the fainter component has been measured.


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There can be no doubt of the reality of its presence. The second spectrum is extremely difficult to measure. When the spectrum of the first component is known, and both minima are thoroughly observed, the second spectrum can be deduced from the difference of surface brightness, obtained from a comparison of the minima. The procedure is as fol- lows: we employ a table of relative surface brightnesses, reduced to the sun as zero. Taking from this table the surface bright- ness of the first component in magnitudes corresponding to the observed spectrum of the first component, we add to it the difference in magnitudes of surface brightness between the two components, as deduced from the light curve; the result is the surface brightness of the second component, expressed in mag- nitudes, and the corresponding spectrum may be read off from the table.

This method has important results, and enables us to compare the observed spectra with the spectra thus obtained from the light curve. In many cases the agreement is good, but in some cases, such as u Herculis, p Lyrae, Z Vulpeculae, RS Vulpeculae, and the recently discussed TX Ursae Majoris, there is a large discrepancy, amounting to one spectral class. It is of some interest to note that the type of eclipse is re- lated to the difference in spectra. RX Geminorum is a strange system with respect to its spectrum; it contradicts all our ex- pectations founded on the period, amplitudes, and the spectrum of the primary.

It is evident that the explanation lies in the relative size of the components; the difference in spectra is closely re- lated to the difference in radii. In almost all cases the difference in spectra is considerable and the spectrum at the first minimum should show special features. The remarks to the table contain descriptions of some of the spectra during eclipse. We give here a few additional details. At the Lick Observatory an investigation has been undertaken by Wyse with the special purpose of determining the nature of the spectra during eclipse for 25 systems.

The star RW Tauri gave remark- able results. At the expected minimum, November 21, , the second spectrum AO showed the bright lines of HI3, Hy, H8, and probably Mgll ; but at the minima of December 16, , and January 10, , there was no trace of the second spectrum or of the bright lines. During the minimum of the remarkable star VV Cephei the bright hydrogen lines disappear, but the bright for- bidden lines of iron do not. Its lines become somewhat stronger, and some display an asymmetry for example, , during the ascending branch of the light curve. Aurigae shows no peculiarities during the eclipse, but at the beginning and end there is an indication of the existence of an atmosphere around the supergiant K5 star.

We may mention here that there appears to be no one-para- meter relation between spectrum and period. Martinoff recently came to the same conclusion statistically. Frost, Struve, and Elvey, Publ.


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  8. Yerkes Obs. RADII 39 tistically, though to a less degree, on the duration of eclipse.

    Variable Stars

    Masses have been determined spectroscopically, with consider- able certainty, for 35 eclipsing systems, but the number of de- terminations can be increased to over , as is explained in the section devoted to masses. Other methods of determining the radius may be mentioned here. Considering the absolute dimensions of eclipsing stars, we note that the largest value is that for the red component of VV Cephei G , and the smallest that for UX Ursae Majoris 0. In the former sys- tem there is a large secondary variation caused by the fluctua- tions of the brighter star. Although UX Ursae Majoris has the shortest known period for an eclipsing star, it is nevertheless described as of Algol type, so that the ellipticity of the com- ponents, if any, must be very small.

    History Topics

    A more detailed study of dimensions leads to the following results: Among the three O stars, AO Cassiopeiae and 29 Canis Majoris have nearly equal radii 23 O , and Y Cygni a much smaller radius about 5 0. One must, however, remember that for the first two the photometric elements are of much lower accuracy than those for Y Cygni. For the other 38 eclipsing stars of Class B we find a mean radius of 4 O. The A stars have a range from 0. Excluding three systems ""Lund Circ.

    Lund Medd.

    The Light Curve of XX Cygni As a Contribution to the Study of Cepheid Variation (1915)(en)(5s) The Light Curve of XX Cygni As a Contribution to the Study of Cepheid Variation (1915)(en)(5s)
    The Light Curve of XX Cygni As a Contribution to the Study of Cepheid Variation (1915)(en)(5s) The Light Curve of XX Cygni As a Contribution to the Study of Cepheid Variation (1915)(en)(5s)
    The Light Curve of XX Cygni As a Contribution to the Study of Cepheid Variation (1915)(en)(5s) The Light Curve of XX Cygni As a Contribution to the Study of Cepheid Variation (1915)(en)(5s)
    The Light Curve of XX Cygni As a Contribution to the Study of Cepheid Variation (1915)(en)(5s) The Light Curve of XX Cygni As a Contribution to the Study of Cepheid Variation (1915)(en)(5s)
    The Light Curve of XX Cygni As a Contribution to the Study of Cepheid Variation (1915)(en)(5s) The Light Curve of XX Cygni As a Contribution to the Study of Cepheid Variation (1915)(en)(5s)
    The Light Curve of XX Cygni As a Contribution to the Study of Cepheid Variation (1915)(en)(5s) The Light Curve of XX Cygni As a Contribution to the Study of Cepheid Variation (1915)(en)(5s)

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