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University Microfilms International 300 N. ZEEB RD.. ANN ARBOR, Ml 48106 8129072

Os w a l t , T e r r y D e a n

A SPECTROSCOPIC STUDY OF COMMON PROPER MOTION BINARIES WHICH CONTAIN DEGENERATE COMPONENTS

The Ohio State University PH.D. 1981

University Microfilms International 300 N. Zeeb Road, Ann Arbor, MI 48106

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University Microfilms International A SPECTROSCOPIC STUDY OF COMMON PROPER MOTION

BINARIES WHICH CONTAIN DEGENERATE COMPONENTS

DISSERTATION

Presented in Partial Fulfillment of the Requirements for the Degree

Doctor of Philosophy in the Graduate School of

The Ohio State University

By

Terry Dean Oswalt, B.A.

*****

The Ohio State University

1981

Reading Committee: Approved By

Professor Terry P. Roark

Professor Gerald H. Newsom

Professor Eugene R. Capriotti

Adviser Department of Astronomy DEDICATION

To my wife, Barbara, whose moral support and (almost)- infinite patience has been greatly appreciated, and to our first child,

Amanda Marie, whose arrival coincided with the frantic final prepar­ ation of this manuscript.

ii ACKNOWLEDGEMENTS

The observational material for this research was obtained at the

Lowell Observatory and the Cerro Tololo Inter-American Observatory.

I wish to thank the staff of these institutions for the technical support necessary for this study. In particular I greatly appreciate the invaluable assistance of John Graham in obtaining CTIO observing time. For advice and encouragement on numerous occasions I thank my advisor Terry Roark and reading committee member Gerald Newsom. I am indebted to Willem Luyten, Henry Giclas, Gary Wegner, and Don Barry who all provided crucial data and/or research results in advance of publication. Computing time was provided by the Ohio State Instruction and Research Computer Center for which I am also grateful. Finally sincere thanks are due Delores Chambers and Kenneth Rumstay who prepared most of the final manuscript and figures, respectively.

iii VITA

October 30, 1952. . . . Born - Wabash, Indiana

1974...... B.A., Indiana University, Bloomingtion, Indiana

1975-1977 ...... Teaching Associate, Department of Physics, The Ohio State University, Columbus, Ohio

1978...... Research Assistant, Perkins Observatory, Delaware, Ohio

1978 ...... Instructor, Department of Astronomy, The Ohio State University, Mansfield, Ohio

1979 ...... Teaching Associate, Department of Astronomy, The Ohio State University, Columbus, Ohio

1979-1981 ...... Instructor, College of Mathematical and Physical Sciences, The Ohio State University, Columbus, Ohio

1981...... Instructor, Department of Astronomy, The Ohio State University, Mansfield, Ohio

PUBLICATIONS

"Current Observations of the White Dwarf Eclipsing Binary V471 Tauri." Bull. Am. Astr. Soc. 10. 410, 1978.

"Call H&K and Hydrogen Line Variations in V471 Tauri (=BD+16°516)." Publ. A.S.P. 91. 222, 1979.

FIELDS OF STUDY

Major Field: Binary Stars (Spectroscopy and Photometry)

Studies in Close Binary Stars. Professor Terry P. Roark

Studies in Common Proper Motion Binaries; Professor Terry P. Roark

iv TABLE OF CONTENTS

Page DEDICATION...... ii

ACKNOWLEDGMENTS ...... iii

VITA...... iv

LIST OF TABLES...... vii

LIST OF FIGURES...... viii

Chapter

I. INTRODUCTION ...... 1

II. THE SOURCE CATALOGS AND SELECTION OF CANDIDATES. 13

Selection of Candidates ...... 17 General Characteristics of the Selected CPM Systems...... 29

III. INSTRUMENTATION AND REDUCTIONS...... 45

Reduction of MC Spectra ...... 49

IV. ABSORPTION LINE MEASUREMENTS IN THE SPECTRA OF DEGENERATE MEMBERS OF CPM BINARIES...... 53

The H-Rich CPM W D ' s ...... 57 The He-Rich CPM WD's...... 74 Conclusion...... 79

V. THE SPECTRA OF NON-DEGENERATE MEMBERS OF CPt; BINARIES...... 80

Quantitative Classification of Late-Type Stellar Spectra: The Cal 4227 and TiO 4954 Indices...... 83 Quantitative Classification of dF and dG Stars: The Call H&K Continuum Index . . . 93 Chromospheric Activity in CPM RD's: Call H&K and HB Emission...... 97 Summary and Conclusions ...... 114

v TABLE OF CONTENTS (Continued)

Page VI. SUMMARY...... 116

APPENDIX

Program VIDIC0N2...... 120

LIST OF REFERENCES...... 135

vi LIST OF TABLES

Table Page

1. Color Characteristics of Luyten Color Classes (Jones, 1972b)...... 24

2. Color Characteristics of Giclas Color Classes (Giclas, et al. , 1971)...... 27

3. Degenerate CPM Stars Observed With the SIT Vidicon. . 56

4. Absorption Line Measurements in the Spectra of DA White Dwarfs...... 63

5. Absorption Line Measurements in the Spectra of DB White Dwarfs...... 77

6 . Non-Degenerate CPM Stars Observed With the SIT Vidicon...... 81

7. Wavelength List for Features Measured in the Spectra of Non-Degenerate CPM Stars ...... 84

8 . Spectral Indices Measured in the Spectra of Non-Degenerate CPM Stars...... 85

9. Spectral Indices Measured in the Spectra of Bright F-M Dwarfs...... 87

vii LIST OF FIGURES

Figure Page

1. The nipg vs. B Calibration for Luyten and Giclas CPM Stars...... 18

2. Comparison of the Luyten and Giclas Color Classes for CPM Stars...... 19

3. The Relation Between Hg, Mg, and T for the Degenerate and Non-Degenerate Members of CPM Binaries...... 23

4. The Effect of Color Class on the mpg vs. V Calibration for Luyten CPM Stars...... 26

5. Hertzsprung Diagram for All Giclas CPM Stars, With WD Candidates Isolated...... 28

6 . Positions on the Celestial Sphere of 208 CPMLIST Binaries...... 30

7. Comparison of the Log y Histograms of (a) Luyten and (b) Giclas CPM Binaries...... 32

8 . Comparison of the Ampg Histograms for Components of (a) Luyten and (b) Giclas CPM Binaries...... 34

9. Comparison of the Ac Histograms for Components of (a) Luyten and (b) Giclas CPM Binaries...... 35

10. Comparison of the Ampg and Ac Values Exhibited by Luyten CPM Binaries, Showing the Empirical Division Between Luyten*s Groups A and B,B'...... 38

11. Hertzsprung Diagram for CPMLIST Binaries Having Published UBV Photometry...... 40

12. Comparison of the AB and |A(B—V)| Values for All CPMLIST Binaries Having Published UBV Photometry, Showing the Empirical Separation of Luyten's Group A and B,B' CPM Pairs...... 42

viii LIST OF FIGURES (Continued)

Figure Page

13. The B' vs. B Calibration for Degenerate and Non-Degenerate Stars Observed With the SIT Vidicon. . 47

14. The Dependence of FWHM(y) and W(y) on Temperature and Gravity in DA WD's...... 60

15. The Dependence of W(6e) and W(8y) on Temperature and Gravity in DA WD's...... 62

16. Comparison of the Instrumental and Theoretical H,He Equivalent Widths for DA and DB WD's Observed With the SIT Vidicon...... 65

17. Comparison of SIT Vidicon and Photographically Determined H,He Equivalent Widths for DA and DB WD's. 66

18. The Effect of Finite Instrumental Resolution on SIT Vidicon FWHM' (y) in DA W D ' s ...... 70

19. The Correlation Between SIT Vidicon C(40-49) and U-V Color for WD's...... 73

20. The Dependence of W(HeI 4026) and W(HeI 4472) on Temperature and Gravity in DB WD's...... 76

21. The Dependence of S(Cal) on Spectral Type in K-M Dwarfs...... 89

22. The Dependence ofS(Cal) on B-V Color...... 91

23. The Dependence of C(HK) on Spectral Type in F-K Dwarfs...... 95

24. The Dependence of C(HK) on B-V Color...... 96

25. Comparison of (a) S(HK) and (b) S(3) Vidicon Indices With Call H&K Fluxes Reported by Wilson (1978). . . . 101

26. The Correlation Between S(3) and S(HK) Among Bright F-M Dwarfs Observed With the SIT Vidicon...... 103

ix LIST OF FIGURES (Continued)

Figure Page

27. The Dependence of S(HK) on B-V Color...... 104

28. The Dependence of S(3) on B-V Color...... 105

29. The Correlation Between AS (3) and AS(HK) Among F-M CPM Dwarfs...... 107

30. Comparison of AS(HK) Exhibited by Components of CPM Binaries Containing Two F-M Dwarfs...... 110

31. Comparison of AS(3) Exhibited by Components of CPM Binaries Containing Two F-M Dwarfs...... Ill

32. Comparison of (a) AS(HK) and (b) AS(3) of F-M CPM Dwarfs to Log T3 of WD Companions...... 113

x CHAPTER I

INTRODUCTION

It is now generally accepted that the white dwarfs (WD's)

represent the most common final phase in the process of stellar

evolution. The discovery that subluminous stars such as Sirius

B, having masses similar to that of the sun yet radii comparable

to that of the earth, led to the conclusion that objects of this

type possess electron degenerate interiors (Eddington, 1924).

Most WD's have been discovered by search programs which

isolate objects of large proper motion and faint apparent mag­

nitude. Usually color criteria are also applied since the

hottest and therefore "bluest" WD's are fairly easy to dis­

criminate from the multitudes of faint late type stars in the

Galaxy that share similar motions and magnitudes. This technique

has worked exceptionally well; for example, Luyten (1977) has

identified about 6500 WD candidates from among some 200 million

stars on his original survey plates.

A second discovery technique utilizes the fact that WD's often exhibit different photometric colors from main sequence

(MS) stars in a color-color diagram. The loci of degenerate objects in such a diagram, e.g. (U-B) vs (B-V), fall generally above the MS line, usually close to colors that would be exhibited

1 by a blackbody. This method is limited by the fact that the hot­

test WD’s having (B-V) < -0.2 are photometrically indistinguishable

from early MS or subdwarf 0 stars while for (B-V) > +0.5 weak-line and late type subdwarfs also lie above the MS curve. It is also difficult to identify spurious objects such as QSO's and composite systems among the faintest objects unless proper motion or infrared photometry is available. Eggen and Greenstein (1965a) and Greenstein

(1969) provide a general description of the photometric UBV pro­ perties of WD's.

The extremely high surface gravities characteristic of WD's are usually manifested by strongly pressure-broadened lines of H,

He, and/or metals. In fact spectroscopic programs conducted by

Greenstein -(1979a, and references cited therein), Liebert and

Strittmatter (1977), and others for WD suspects identified by proper motion surveys have relied primarily on the presence of such spectral features. Faint, high proper motion stars with nearly featureless spectra are also often classified as WD's.

Early attempts at spectral classification by Luyten (1945) added a prefix "D" (indicating a degenerate object) to spectral types similar to those for MS objects if the WD spectrum somewhat re­ sembled or had similar colors to that of a MS star. Additional designations such as DC and A4670 (now generally called C^, cf.

Wegner, 1973) were reserved for WD's exhibiting continuous spectra and features now ascribable to the absorption bands of

C£. WD's normally display absorption lines of only one element in a single stage of ionization. The classification of de­ generate star spectra is therefore now usually established by criteria similar to those used by Hintzen (1975), the most common­

ly occurring types being DA (hydrogen lines), DB (helium lines),

and DC (continuous spectrum).

The WD's for which UBV photometry (Johnson, 1952) and re­ liable parallaxes are available define a sequence in an H-R diagram extending roughly between My = +8 , (B-V) = -0.4 and

My = +16, (B-V) = +1.1 which is approximately parallel to and

5-10 magnitudes fainter than the MS (see Harrington et al., 1978).

Some indication of a general convergence of these two sequences exists towards larger (B-V) colors, due primarily to the downturn in the MS towards later spectral types.

To avoid the effects of broad absorption lines on the B magnitudes Eggen and Greenstein (1965a) chose (U-V) as the color index most useful for determination of the luminosity-temperature relation for WD's. Unexpectedly, two sequences were apparent in

My, (U-V) plots for WD's. Graham (1972) showed this bifurcation to be a consequence of line blocking (particularly in DA WD's) which affects the interpretation of UBV photometry in a very complicated way. A single tight sequence emerged when the uvby photometric system (Stromgren, 1963) is used to construct an

My, (b-y) diagram. In addition the v filter was demonstrated to be particularly sensitive to the presence of H6 absorption. Un­ fortunately this intermediate band system is difficult to apply to the study of objects fainter than V = +16 with telescopes of moderate aperture. The WD sequence has generally been Identified as an evolutionary

sequence for objects which derive their emitted energy from the

cooling of electron degenerate interiors, rather than nuclear re­

actions (Mestel and Ruderman, 1967). Thus it may also be properly

termed a cooling sequence, and the locus of a particular WD in an

H-R diagram contains information about the mass and duration of time

the object has been losing thermal energy.

The Mestel and Ruderman theory predicts substantial numbers

of observable cool WD's in the solar neighborhood. Calculations

imply a gradual decrease in the cooling rate of an electron de­

generate core as its temperature decreases until the Debye tem­

perature is reached and/or the ions become degenerate. These

changes in specific heat of the core allow the star to rapidly cool

although, as before, the actual rate of energy loss can be appreci­

ably altered by the opacity of the nondegenerate atmosphere.

Observationally these objects fade relatively quickly once they

reach M^ = +16. Lamb and van Horn (1975) have discussed the im­

plications such a cooling theory has for observational searches for

cool degenerates.

The results of early spectral surveys (e.g. Greenstein, 1969)

suggested that far fewer cool degenerates having (B-V) > +0.4 than

predicted by cooling theory were being found. The difficulties associated with the identification of cool degenerate stars stem

from their intrinsically low luminosities (M^ > +14) and the con­ fusion generated by overwhelming numbers of MS stars and high- velocity subdwarfs which exhibit similar proper motions and colors. The first of these two problems has been attacked with

limited success by Hintzen and Strittmatter (1974), Liebert (1978),

and Hintzen (1979) by searching to fainter limiting magnitudes

than that achieved by previous investigations (e.g. Greenstein,

1971). Although these studies significantly increased the number

of known cool degenerates with < +15 the detection rate at ’

fainter luminosities was disappointingly low. Only within the

severely limited volume sample of stars nearer than ^5 pc has the

number of known cool WD's been shown to be consistent with that

indicated by theoretical cooling times (Weidemann 1967; Green,

1977; and Sion and Liebert, 1977). This implies that detection

rates beyond a few parsecs from the sun remain poor. The situa­

tion for the coolest degenerates (M^ > +15) is even worse due to

the observational difficulties such faint objects pose for spec­

troscopic work and their presumed scarcity.

Using computed line profiles for H, Call H&K, and Cal 4226

Hintzen and Strittmatter (1975) demonstrated that cool hydrogen-

rich degenerates could be expected to show spectra not easily

distinguishable from subdwarfs at moderate spectral resolutions.

More recently, however, Wickramasinghe et al. (1977) found that

WD's as cool as 5000K would exhibit Balmer-line spectra much weaker than subdwarfs of the same effective temperature, while metal lines such as Cal 4227 would appear significantly stronger.

These results were credited to a more realistic treatment of con­ vection in the atmospheric models and were shown to be readily detectable at spectral resolutions of ^R. Some of the spectra of

cool degenerates exhibited by Liebert (1979) qualitatively support

this idea though the matter is still a subject of debate.

At still cooler temperatures (B-V) > +1.0, and the sensi­

tivity of the strengths and shapes of molecular bands due to dif­ ferences in surface gravity in nondegenerate stars is well known.

This leads to the conclusion that it is difficult, if not im­ possible, for the coolest degenerate atmospheres to produce spectra that mimic the subdwarfs. It was on this basis that

Liebert (1975) eliminated the only remaining "DM" candidates, since red spectra of these objects could be shown to be similar to those of late type non-degenerates.

Presumably stars with initial masses up to several times the

Chandrasekhar limit (Chandrasekhar, 1939) may become degenerate objects if various post-MS mass loss mechanisms discussed in the literature occur (see for example, Savedoff, 1966). This is particularly evident since it is known that relatively young clusters such as the Pleiades and Hyades which contain WD's have

MS "turnoff" masses in excess of 1.4 Mq (Tinsley, 1974). Cluster membership provides an independent age estimate which, if a par­ ticular cooling theory is assumed, can be used to derive WD progenitor masses (Sweeney, 1976). Unfortunately no such in­ dependent timescales can be obtained for single WD's.

Binary stars have provided exceedingly well-used methods for the determination and calibration of stellar astrophysical pro­ perties. While such information is commonly assumed to also apply

to single stars this assumption may be vitiated if binary system

membership has significantly influenced evolution of the component

stars at some past epoch (see for example, Batten, 1973).

Within the last two decades it has become evident that re­

latively wide common proper motion (CPM) pairs constitute the most

frequent type of binary systems (Luyten, 1971). The proper motion

catalogues of Luyten (1963), Giclas et al. (1971, 1978), and others

contain thousands of such systems which usually consist of two

late-type dwarfs. Approximately 40 per cent of the red dwarf stars

in the McCormick and Michigan objective prism surveys have been found to be members of CPM pairs with separations of a few hundred

AU or less (Worley, 1969), implying orbital periods on the order 3 4 of 10 -10 years. An early statistical investigation of CPM pairs found in the course of the Bruce Proper Motion Survey (Luyten, 4 6 1933) provided period estimates in the range 10 -10 years, im- 3 4 plying separations of 10 -10 AU. Indeed, it has been shown that moving pairs having spatial separations as large as 10-20 pc exist 9 (suggesting orbital periods ^10 years!) in numbers far higher than that which would result from a random distribution of space velocities (Lii and Upgren, 1973). Taking these results into account, about 60 per cent of the McCormick and Michigan survey stars having well-determined space velocities are members of very wide pairs. This proportion is nearly independent of velocity, implying relaxation times which are long compared to the age of the Galaxy. Unlike earlier spectral types, dK and dM stars as a group present a wide range of ages.

There are now several hundred CPM pairs which, on the basis of proper motion and estimated colors, are expected to contain a degenerate member (see for example, Luyten, 1979; Giclas, et al.,

1971, and Gliese, 1971). Statistically the companion to these stars is likely to be a red dwarf or subdwarf.

The large physical separations inferred for CPM pairs pre­ clude any quick orbital determinations; consequently little pro­ gress beyond the preliminary work done by Luyten (1961) has been made along these lines. Astrometric programs for a few CPM pairs showing detectable orbital motion are being conducted (see

Harrington et al., 1978 and references therein) but even the most favorable systems will require many years of observation before dynamical masses may be computed. In principle considerable in­ formation about a degenerate component beyond that commonly obtain­ ed for a single object can still be obtained if the system can be assumed to be a physical pair. This appears to be a valid assump­ tion for those pairs exhibiting relatively large annual motion

(y > 0 .2"/y) since the likelihood of a random distribution of space velocities giving rise to such a coincidence of position angle and motion as well as apparent proximity is demonstrably minute (Lli and Upgren, 1973).

Additional information can be derived if it is assumed that members of a given pair are coeval. Observational evidence con­ tinues to strengthen the hypothesis that stars are born in 3 8 unstable groups which disintegrate in 10 -10 years leaving wide

binary and multiple systems (Batten, 1973). Eggen (1979b) contends

that components of CPM pairs might be former cluster members which

do not necessarily share a common age, since pre-MS late type dwarfs

and the progenitors of WD's in clusters (e.g. the Pleiades) must

have very different contraction times (Hayashi, 1961). Wilson (1963)

however, has shown Call H&K emission strengths among members of a

particular cluster or visual binary system are comparable. This

finding led to the realization that H&K emission (HKe) is a strong

function of stellar age (Wilson and Skumanich, 1964; Wilson 1968).

These results, coupled with the dynamical problems involved with the

formation of captured pairs, imply that the coeval assumption may

still be regarded as a fairly strong one (cf. Batten, 1973, Ch. 10),

although differences in pre-MS contraction times between components

might be expected to impose secondary observational effects.

Eggen and Greenstein (1965a, 1965b, 1967) and Greenstein

(1969) used the MS component in several CPM pairs to obtain lumin­

osity estimates for the WD companions; however aside from work on

five wide binaries by Wegner (1973), no astrophysical data are

available for the vast majority of such systems. Consequently the wealth of information that such physical pairs can provide has

remained largely untapped. Present theories on degenerate stars

rest upon an uncomfortably restricted empirical basis since, for

example, even reasonably reliable mass determinations exist for 10

less than a half dozen WD's which are members of close binary systems (Shipman and Sass, 1980).

It is proposed in this dissertation that CPM pairs provide an avenue through which the empirical data on WD's may be broad­ ened since the complications imposed by uncertain evolutionary influences (e.g. mass exchange in close pairs) may be largely avoided while simultaneously retaining many of the observational advantages afforded by binary systems. With this in mind a large group of faint (m < +17) CPM pairs consisting of a known or P 8 suspected WD and a late type nondegenerate star has been selected from the catalogs of Luyten and Giclas. A discussion of the se­ lection criteria applied and a comparison of the basic kinematic, photometric, and colorimetric properties of this observational sample is presented in Chapter 2.

The present study has gradually evolved into an investi­ gation of the spectrophotometric properties of faint CPM systems containing WD's. Such a program has not been possible with tele­ scopes of moderate aperture until the relatively recent advent of high quantum efficiency, linear response, multi-channel detectors such as the CTIO SIT vidicon. Spectra of nearly 50 systems were obtained with this instrument mounted on the CTIO 1.5m telescope.

The operational characteristics of this system and reduction pro­ cedures have been discussed in detail by Atwood et al. (1979) and are summarized in Chapter 3. It has been demonstrated (cf. Greenstein, 1979b; and Liebert,

1975) that digital spectra obtained with modern multi-channel detec­ tors are readily usable for accurate quantitative measures of atomic lines and molecular band strengths. Relatively little work however has been published on line strengths in faint WD's, and a major portion of Chapter 4 is devoted to a discussion and comparison of the SIT vidicon data with the results of previous photographic spectroscopic surveys. Particular attention is given to the cali­ bration of temperatures and equivalent widths of spectral features common to WD's.

Any investigation of such an exploratory nature is certain to uncover objects of particular individual interest. Table 3 in

Chapter IV lists roughly one dozen WD candidates which have now been spectroscopically identified as degenerate stars. Most of these were originally listed as possible WD's by Luyten (1979). Several peculiar systems are also noted which are especially worthy of additional attention.

Wilson (1963) has shown the strengths of HKe in late type dwarfs to be a strong function of stellar age. Relative measures of HKe strengths in non-degenerate components of the CPM pairs in the present study therefore place age constraints on the degenerate companions that provide an independent empirical limit on theoret­ ical WD cooling times. Especially pertinent to this portion of the study is the establishment of suitable temperature criteria for the late type dwarfs. A discussion of the historical background and observational results of this study appears in Chapter 5. 12

Chapter 6 presents a summary of these investigations. An assessment of the results and comparison with other work that has been done on such pairs is made. Particularly promising topics deserving further work are examined.

The proper motion catalogs contained the only data available for most of the objects in this study. Consequently a large portion of time was spent constructing Palomar Observatory Sky Survey (POSS) finding charts for the 208 systems which were ultimately selected.

Any of these are available from the author upon request. CHAPTER II

THE SOURCE CATALOGS AND SELECTION OF CANDIDATES

In order to establish the general spectroscopic properties of wide binaries which contain degenerate components the present study required a large sample of such pairs. It was also desirable that

the selected systems be well-distributed over the sky to insure

that a sufficient number of spectra could be obtained regardless of the allocated observing time. Unfortunately only a small fraction of the WD's for which spectroscopic and/or photometric identifications exist are also members of CPM binaries (cf. Eggen and Greenstein, 1965a,

1965b; Wegner, 1975; and Greenstein, 1976). The Proper Motion Survey with the Forty-Eight Inch Schmidt Telescope (Luyten, 1963) and the

Lowell Proper Motion Survey (Giclas et al., 1971) provide apparent photographic magnitudes and color estimates in addition to proper motions for large numbers of stars which, in the absence of independent photometric data of spectra, are especially useful for identification of possible degenerate stars (cf. Liebert, 1979 and references therein).

Both surveys provide fairly uniform coverage of the northern sky and significant extensions to the southern hemisphere. For these reasons they were used as the basis for candidate selection. The choice of two independent sources is justified below by the complementary observational advantages each catalog provides.

13 14 In this chapter we present a general description and com­

parison of the Luyten (LP) and Giclas (G) catalogs, followed by a

discussion of the selection criteria which have been applied to

obtain the present observational sample. Even the crude photo­

metric and colorimetric data provided by the source catalogs allow

some general properties of the selected CPM pairs to be distin­

guished. Correlations between magnitudes and colors of companions

have been noticed previously by Luyten (1979) and are apparent in

the present sample. In the closing section of this chapter these

relationships and their probable origin are examined.

The Proper Motion Survey with the Forty-Eight Inch Schmidt

Telescope (Luyten, 1963) covers the area north of declination

-33° included in the original Palomar Sky Survey. This program

has provided the largest and most complete compilation of faint

proper motion objects and will be particularly useful in the de­

termination of the luminosity function for cool WD's (Liebert,

1979) and the search for nearby faint red dwarfs (Gliese and

Jahreiss, 1980). With the exception of the dense regions near the

galactic equator, Luyten (1970) has measured the original Palomar

Schmidt plates and second epoch red plates for all stars with

y > 0.1 "/y down to near the plate limits of m M-21.2 and PS M-19.2. Magnitude estimates based on a single eye interpolation

between the known magnitudes of BD stars on the plates, or in some

cases determined by the Luyten-Control Data Stellar Proper Motion

Measuring Machine (Luyten, 1970), and color classes b, a, f, g, k, 15 m, m+, correlated with the color index ci^ = m - m^, are given for

each object in the catalog. The relation between Luyten's color

classes and B-V color index will be discussed later (see Table 1).

The LP catalog contains a much larger sample of CPM pairs

than any other survey largely because of its lower proper motion cutoff and it also provides more extensive coverage of the south­ ern hemisphere than the G catalog, a fact of particular importance to the present investigation since most of the observational material for the present study was obtained at the Cerro Tololo

Inter-American Observatory. Furthermore Luyten (1969, 1974, 1979) has drawn special attention to CPM pairs containing suspected WD's; over 400 such systems have now been published. Probably because finding charts are largely unavailable these objects have received very little attention from observers; hence the opportunity for spectroscopic identification of many suspected WD's made the in­ clusion of these syterns desirable.

The Lowell Proper Motion Survey (Giclas et al., 1971) is now com­ plete for the northern hemisphere and includes data for all stars de­ tected with p > 0.27 "/y and brighter than m M-17. In addition PS to m and color estimates -1, 0, ... +4 based on comparison of red-blue plates taken with the 13-inch photographic telescope at

Lowell, excellent finding charts are provided for each catalog object. Probably for this reason, extensive photoelectric and spectroscopic data exist for these stars. This survey has been heavily exploited in the search for degenerate objects and other low luminosity stars (cf. Giclas, 1969). The opportunity for calibration and/or comparison of the observations reported here

with previous work therefore provided a strong incentive for the

inclusion of the Giclas CPM pairs in this study.

There exist many Giclas objects which have not been included

in his regular catalog because they exhibit proper motions lower

than the usual survey limit. Because the LP- and G-catalogs differ

’"greatly in motion limit (y > 0.10 " /y and y > 0.27 " /y respectively)

such low y Giclas CPM pairs have been considered as supplementary

program objects. The sample of Giclas pairs selected for the pre­

sent study therefore contains three distinct groups which differ

in the selection criteria applied. First, the northern hemisphere

catalog and its southern extension (Giclas et al., 1978) were searched

for CPM pai.rs containing at least one WD identified by independent

photometry or spectroscopy. For this group y > 0.27 " /y. Second,

Giclas has published occasional lists of possible pairs exhibiting

motions below his regular catalog limit (hence y < 0.27 " /y) which

have contrasting colors. These lists, scattered throughout Lowell

Bulletins #120-163, are indicated by a B prefix to the object number

and are denoted as GB objects henceforth. Many of these objects

have preliminary UBV photometry by Eggen. It should be noted that

physical association of the components of the GB pairs remains

subject to considerable uncertainty. The third group consists of

pairs composed of late type components of color class +2, +3, or

+4 screened from a compilation of all CPM pairs in the Giclas

catalog. These reddish stars will be referred to as GR objects for

convenience. Most have no available photometric or spectroscopic observations and were selected on the basis of large proper motion and color class criteria chosen to isolate prospective cool de­ generates. These criteria are discussed in the next section.

To facilitate the estimation of exposure times for objects lacking photoelectric data it was desirable to determine the re­ lationship between the magnitude scales and color classes of the

Luyten and Giclas catalogs. Figure 1 presents the (m , B) Pg calibrations for the Luyten (L) and Giclas (G) CPM objects having published UBV photometry. Evidently a zero point difference of

'V' 0.7 magnitude exists between the two catalogs. No color de­ pendences are evident in the calibrations since the effective wavelengths of B and m are similar. The relation between color Pg classes of objects listed by both catalogs is displayed in Fig­ ure 2. The correlation is rough at best and deviations of +1 color class are common.

Selection of Candidates

Together the G, GB, and GR objects span roughly the same range in p as the Luyten sample though, as mentioned earlier, different selection criteria have been used to define each group.

All the selection criteria however must discriminate stars of the desired intrinsic luminosity and temperature, i.e. the objects on the WD sequence in a Hertzsprung-Russell (H-R) diagram.

Three fundamental quantities are required to establish the true locus of a particular star in an H-R diagram: its apparent magnitude, spectral type (or color), and parallax. Because the 15 m pg

10

• LlfYTEM + GICLAS

10 15 20 B

Figure 1. The m vs. B Calibration for Luyten and Giclas CPM Stars. Pg

00 iue2 Comparisonofthe andLuytenColor Giclas Classes for Figure2.

COLOR CLASS (LUYTEN) b« ob km • • I 9k IS o b k CPM CPM Stars. 2 1 OO LS (GICLAS) CLASS COLOR 0 +1 ■H-

+2

+3

+4 19 20 parallax Is difficult to obtain for large numbers of faint stars a

statistical correlation between stellar proper motion and parallax

often provides the only data.on the intrinsic luminosities of such

stars. One formulation of this concept has been extensively em­

ployed by Luyten (1977) who credited its origin to Hertzsprung.

For this reason the "reduced proper motion" was symbolized by H

and is defined by

H = m + 5 log p + 5 II-l where m is apparent magnitude and p the total annual proper motion

in seconds of arc per year. The transverse velocity of a star T, in AU per year, is related to its observed proper motion and parallax it, in seconds of arc, by the small angle approximation

T = p/ir . II-2

Since distance modulus is defined as

m-M = 5 log - - 5 II-3 7T equations II-2 and II-3 can be used to relate a star's reduced proper motion to its absolute magnitude M and transverse velocity

H = M + 5 log T . II-4

Therefore, at a given color, H reflects the combined absolute magnitude and velocity distributions of stars in the solar neighbor­ hood. 21

Because the dispersion in T within each of the two major

stellar velocity groups (Population I and Population II) is gen­

erally smaller than the dispersion in intrinsic luminosities, H

can replace M in a statistical analog to the H-R diagram (see for

example Jones (1972a) and Luyten (1977) ). Such a plot is called a Hertzsprung diagram and will be of use in the following dis­ cussion.

The photographic magnitudes and colorimetric data provided by the Luyten and Giclas catalogs are particularly useful for the identification of possible degenerate members of CPM binaries since the reduced proper motion, relative luminosity, and colors of the components may be compared to those of likely pairings in a

Hertzsprung diagram which has been calibrated by a well-determined

H-R diagram. This procedure requires determination of the re­ lationships between the corresponding axes of each plot.

The relation between H and M may be obtained from Equation

II-4, by adopting a mean transverse velocity for the group of objects sought. For example, Jones (1972a) identified the giant branch, main sequence, subdwarfs, and WD sequence in an H^ vs. B-V

Hertzsprung diagram for nearly 8000 objects by assuming = 100

AU/y for subdwarfs and = 6.33 AU/y for all other stars. On the other hand an empirical relation between H and M may be derived from proper motion objects of known intrinsic luminosity. Luyten

(1952) found M MD.90 H -3.3 for a group of WD's having Pg Pg 6 parallaxes determined trigonometrically or from membership in a cluster or binary system. Since some known WD members of Luyten 22 and Giclas CPM pairs now have reliable trigonometric parallaxes and

UBV photometry (Harrington et al., 1978) a similar relation can be

established for these objects. Figure 3 presents the correlation

between Hg and Mg for these systems. From Equation II-4 lines of

unit slope should characterize stars of equal T. The transverse

velocities adopted by Jones (1972a) for subdwarfs and Population I

objects are indicated in the figure. The. points A, B, and C rep­

resent the positions of the respective components of the well

known high velocity triple systems o^ Eri (cf. Gliese, 1971).

Evidently H runs about five magnitudes larger than M,. implying o D that most of these pairs exhibit transverse velocities similar to

Eri, i.e. characteristic of the disk or intermediate Population

II stars. This appears to contradict the results of Iwanowska

(1973) and Sion and Liebert (1977), who found that nearly all

single WD's kinematically appear to be relatively low velocity old

Population I objects. It should be noted however that parallax

programs of faint stars preferentially include objects of large

proper motion. Because of this same bias the absence of WD's with

transverse velocities approaching 100 AU/y implies that very few

degenerates exist among the subdwarfs (cf. Greenstein 1965).

Hertzsprung diagrams are most useful when the relation be­

tween color class and photometric color is known. For example,

Jones (1972b) showed the degenerate sequence is clearly defined in

plots of Hy vs. (B-V). In conjunction with the reduced proper motion - absolute magnitude calibration previously discussed 23

24 ^ '

22

20 Hb

• WD + RD

8 10 12 14 16 M b Figure 3. The Relation Between Hg, and T for the Degenerate and Non-Degenerate Members of CPM Binaries. criteria can then be established for objects lacking photometric colors which will be most likely to discriminate degenerate stars in a reduced proper motion vs. color class diagram.

Table 1. Color Characteristics of Luyten Color Classes (Jones, 1972b). < H > Luyten Color Class n a pg

b 4 0.07 0.13 13.5

a 97 0.14 0.21 14.4 a-f 10 0.50 0.21 - f 61 0.56 0.28 15.5 f-g 6 0.67 0.21 - g 136 0.83 0.36 17.0 g-k 3 1.41 0.21 17.7 k • 133 1.24 0.33 18.3 k-m 11 1.31 0.23 - m 47 1.33 0.20 -

The mean (B-V) colors for samples of various Luyten color classes were published by Jones (1972b) and are reproduced in

Table 1 along with Luyten’s (1979) assumed lower limits-on H pg for degenerate objects. Considerable dispersion in (B-V) exists for stars of a given color class. The coarse Luyten color classes b ... g are sufficiently separated in color to discriminate the hottest degenerates. Unfortunately color classes k ... m are hopelessly intermingled and would therefore be expected to yield large numbers of subdwarfs among the cooler WD candidates. Jones attempted to circumvent this problem by adopting more stringent limits on H at each WD color class Jthan those used by Luyten. 25 While objects selected by Jones' criteria are almost certainly WD's

an appreciable number would be overlooked because he assumed a

negligible color dependence in the (m , V) calibration. That a PS color class dependence is not negligible is evident in Figure 4.

Correction for this effect by adoption of a color class term in the

(m , V) relation, or the more straightforward calibration between Pg m and B magnitude (see Figure 1) results in criteria more P S closely in line with those of Luyten (1979). For this reason

Luyten's criteria have been implicitly adopted by the selection of

128 CPM binaries from his selected lists (1969, 1974, 1979) which have components brighter than m M-17. In retrospect the high yield of new spectroscopic identifications of WD's among this list during the course of the present study (see Table 3) is testimony to the validity of Luyten's selection criteria.

The mean (B-V) colors of the five Giclas color classes are reproduced in Table 2 (cf. Giclas et al., 1971). As is the case for the Luyten objects, color classes are only loosely correlated with photometric colors. Individual members of a class may easily differ in color from the group mean (B-V) by a half magnitude or more. Even so, stars of color class -1 or 0 have been shown to be mostly WD's

(Eggen, 1968). Hintzen (1975) and others found significant numbers of yellow degenerates among stars of color class +1. For redder color classes confusion between the coolest degenerates and sub­ dwarfs results from their similar magnitudes and colors coupled with the relatively bright limiting magnitude of the survey 26

20 COLOR CLASS O a .a f

▲ S.0k x k,km + m,m+

la) 15

Vs A X A •

10 (O)

5

5 10 15 20

Figure 4. The Effect of Color Class on the mpg vs. V Calibration for Luyten CPM Stars. 27

Table 2. Color Characteristics of Giclas Color Classes (Giclas et al., 1971). Giclas Color Class < H > Pg

-1 +0.11 15.1

0 +0.26 16.3

+1 +0.80 19.0

+2 +1.30 20.0

+3, +4 +1.61 20.6

(nipg ^ +17). The GR pairs, having reddish components of color class

+2, +3, or +4, are subject to this difficulty and were the only

Giclas CPM systems in the present study which did not contain in­

dependently identified WD's. The selection process for these

systems will now be discussed.

Figure 5 presents the Hertzsprung diagram for all Giclas

CPM pairs using the mean (B-V) colors of the five color classes as

the abscissa since the latter do not represent equal steps in actual

photometric color (see Table 2). Color class +4 has few. members

and those having photoelectric photometry appear to have (B-V)

colors similar to +3 objects. For this reason these two groups will be combined henceforth.

The My, (B-V) data for WD's having reliable parallaxes and UBV photometry given by Harrington et al. (1978) outline a degenerate

sequence roughly composed of two adjoining line segments. WD's of

(B-V) < +0.5 generally lie along the line My =7.4 (B-V) + 10.8 1---- 10 KNOWN WDs o GICLAS CPM + 11 12 13 14 15 16 17 18 19 20 21 22 23 COLOR CLASS 24 -1 O +1 +JB +3 . w , I . f 0.0 1.0 2.0 B-V i. i srtzsprung Diagram for All Giclas CPM Stars, Wj ) Candidates Isolated. while the more rare degenerates having (B-V) > +0.5 indicate a

sequence of shallower slope, My = 1.2 (B-V) +13.9. The scatter in

My about either sequence is about +0.8 magnitudes. Assuming these

two degenerate sequences and Mg = My + for each Giclas color

class allows Hg to be calculated if T =6.33 AU/y is representative

of old population 1 stars. The small correction for the relation

between B and m given in Figure 1, which also applies to the PS relation between Hg and H t allows the WD sequence to be plotted

in Figure 5. This appears as the two lower line segments in

the diagram. The criteria actually used are the two line segments

offset upwards by ^0.8 magnitude to allow for the observed dis­

persion in My along the WD sequence. The single WD's of color

class 0 and- +1 spectroscopically identified by Hintzen (1975) are

plotted with special symbols in the figure. Clearly the above

criteria work very well for color classes earlier than +2. Objects

falling below the line H = 17.4 + 2.0 for color classes >+2

are components of the 12 GR pairs in this study. It is readily

seen that the Giclas catalog is unlikely to produce many cool de­

generates among the CPM pairs.

General Characteristics of the Selected CPM Systems

Presented in Figure 6 are the positions on the celestial sphere of 208 Luyten and Giclas CPM pairs selected by

the methods discussed in the previous section. Those pairs for which spectroscopic observations were obtained during this program are denoted by special symbols. The early southern hemisphere ••I

•UNOBSERVED • IMAGE TUBE 0SIT VIOICON « I T A SIT

Figure 6. Positions on the Celestial Sphere of 208 CPMLIST Binaries.

u> o 31

Bruce Proper Motion Survey (Luyten, 1941) has yet to be superseded by programs which approach the completeness of the northern hemi­ sphere Palomar Schmidt survey and it is apparent that many faint southern CPM systems remain undetected.

All published information and cross-references found in the literature for each of the selected CPM systems have been collected

Into a punched card compilation named CPMLIST. Software exists that allows retrieval of any item contained on the cards and current epoch positions, offsets, and finding charts can be generated. This collection of data allows an examination of some general properties of the CPMLIST objects. We first discuss the only information available for all CPMLIST stars: the proper motions, photographic magnitudes, and color estimates provided by the Luyten and Giclas catalogs.

Frequency distributions for the motion data are presented in

Figure 7. The log p histogram for Luyten objects (Figure 7a) suggests an increase in numbers of WD pairs that is proportional to -1.2 -3 p ' , rather than to p which would be expected from Equation II-2 and the increasing space volume implied by smaller values of p. A similar histogram for the Giclas objects (Figure 7b) is markedly affected by the inclusion of the GB pairs of small motion

(log p < -1.0). The second maximum at -0.6 < log p < -0.4 may be ascribed to objects near the regular Giclas catalog motion limit, and the third hump at -0.2 < log p < +0.2 appears to be due to the selected GR pairs of large proper motion. For these reasons a com­ parison of the Luyten and Giclas motion data cannot be made. OGO«^NON^«fiOO *•••*<*»**«log 0 /A m V I I I I I I I I I OOOvO'J’CMOCSSfvD ..«.*«««« tHOOOOOOOOI I I I I * + +

Figure 7. Comparison of the Log y Histograms of (a) Luyten and (b) Giclas CPM Binaries. 33

Figures 8 and 9 exhibit histograms for the magnitude

differences Am and color class differences Ac, respectively. All

differences are expressed in the sense secondary minus primary. The

primary is defined as the photographically brighter component and

the discrete color classes b through m or -1 through +4 are treated

as integral steps. Color differences involving intermediate classes

such as k-m, etc. are denoted by an h suffix indicating half-color-

class steps.

A m histograms for the Luyten and Giclas pairs (Figures 8a and ro 8b, respectively) appear similar and emphasize the predominance

of CPM pairs having small photographic magnitude differences. A

second maximum (labelled B) exists between Am = 3.0 - 3.9 in Pg the Luyten sample. The Giclas pairs, (Figure 8b), on the other

hand exhibit a relatively smooth distribution in A m with no Pg secondary maxima.

The diagrams for Ac are most interesting. Both the Luyten and

Giclas Ac histograms (Figures 9a and 9b, respectively) appear multimodal. The overall similarity of these two diagrams (i.e. the

preponderance of pairs having A c < 0 and the second maximum at

Ac > +3) is surprising since the Giclas objects do not represent a uniformly-selected sample and, more importantly, because the physical association of the GB pairs is uncertain. In addition the corre­

spondence between the Luyten and Giclas color classes is rough at best (see Figure 2). 34

-- ®fHc^cn

Figure 8. Comparison of the Ampg Histograms for Components of (a) Luyten and (b) Giclas CPM Binaries. a x. x. x x x A O srfp«NrHXi-ICMn»a- LA V I I I I + + + + + v I I I I I I I A Cl N rt i r( N (O I I I I + + +

Figure 9 Comparison of the Ac Histograms for Components of (a) Luyten and (b) Giclas CPM Binaries. 36

Most of the fine structure apparent in the Am and Ac histo- Pg grams can be explained by the existence of three categories of CPM

binaries containing suspected WD's originally identified by Luyten

(1974) in Hertzsprung diagrams for the LP objects:

(A.) Pairs which seem to be composed of one reddish main

sequence star and a white dwarf with color class f or earlier.

(B.) Pairs composed of one main sequence star and one yellow degenerate with color f-g or later.

(C.) Pairs which seem to be composed of two degenerate stars.

Among the third group of CPM pairs it was found that generally the fainter component was of later color class (i.e. redder) as would be expected if such systems consist of stars on the same degenerate cooling sequence. Photographic and colorimetric discrimination be­ tween hot WD - red dwarf pairs and double degenerates is straight­ forward even with the low accuracy magnitudes and colors provided by the LP catalog; however Luyten was unable to account for the empir­ ical distinction found between WD-main sequence pairs of group A and B.

Luyten (1979) gave mean values of H and color class for Pg components of groups A and B. These have been used to compute mean values of Am (= A H ) and Ac whose approximate positions are marked Pg Pg in Figures 8a and 9a for the LP objects selected for the present study. In each case the correspondence between or < Ac> for one of Luyten's groups and a local maximum in the histogram is quite good. 37 The major features of the Am histograms for both the Luyten and

Giclas CPM pairs appear reasonably well explained. Note the symmetrical maxima at |Ac| > +4 for objects of group A in Fig­ ure 9a. This results from the definition given A c and the fact that such systems have components of similar magnitude, hence in a par­ ticular case either the blue or red component may appear slightly brighter. Only the local maximum at -2h < Ac < -2 remains un­ accounted for in Figure 9a. It is not clear whether these systems should be considered as part of group A; however the deep minimum at the corresponding +2 < A c < +2h implies that it should be re­ garded as an independent feature of the plot (this point will be examined again shortly). It appears that the color differences of groups A and B are indistinguishable for the Giclas objects of Figure 9b although pairs of group A with Ac > +3 seem evident.

Figure 10 summarizes the magnitude and color differences of the LP sample selected for the present study. A noteworthy feature of this diagram is the apparent gap along a line roughly coinciding with Am = 0.55 I A cl +0.6 in an otherwise random dis- Pg 1 1 tribution. The mean values of ^ m ^ an(^ | A c | for group A and group

B pairs are marked. A similar plot for the Giclas objects re­ vealed no analogous gap even when the GB pairs of uncertain CPM are discounted due to the previously mentioned poor separation between groups A and B in the A c histogram (Figure 9b).

The existence of two distinct groups among CPM pairs con­ taining a WD and late type dwarf component is difficult to explain if, as Luyten (1979) has assumed in deriving selection criteria for 10 -

Am pg

5 -

4 o -

Figure 10. Comparison of the Ampg and Ac Values Exhibited by Luyten CPM Binaries, Showing the Empirical Division Between Luyten’s Groups A and B,B'. 39

WD's, a single degexierate sequence roughly parallels the main sequence. A more detailed examination of these two groups is possible using UBV data which are now available for some of these systems. Figure 11 presents the H„D vs. (B-V) diagram for all CPMLIST systems for which UBV photometry of both components is available.

The majority are in fact Giclas pairs. This figure bears a close resemblance to the Hertzsprung diagrams used by Luyten; however here the WD's and late type dwarfs are much more clearly separated.

For comparison, dashed lines representing the main sequence and WD sequence have been drawn using My, (B-V) data from Allen (1973) and the relation between Hg and Mg derived from Figure 3 for low (I) and high (II) velocity objects. Using this relation, the relation between m and B provided by Figure 1, and the mean (B-V) colors ro of Luyten's color classes in Table 1, approximate positions for components of group A and B binaries are plotted in Figure 11.

In addition the mean (B-V) of Luyten's color class f which marks the division between WD components of groups A and B is marked.

The source of Luyten's groups A and B appears to be a conse­ quence of the discontinuities near HcD = +20, (B-V) = +0.5 on the WD sequence and Hg = +14, (B-V) = +1.3 for main sequence stars.

Note that group A pairs can be easily located as WD's of (B-V) <_

+0.5 combined with a late type dwarf of (B-V) +1.3. Since for

Hg j> +14 the WD and main sequences are almost parallel the resul­ tant pairs have similar Hg (or H ) an^ widely different colors. 40

T

O WD 6 • • • RD MSI N CPM: 8 n •V 10 M S I .% lsv W - v HB I2 / V “V 14 ; \ * ,

16 x * &, ° • •. y «!T 18 . _•

20 0 > t° °0 V '“'W D *I O •*& • 22 o s v ° ° 24 ' w o n

j_____ j_ 0.0 + 0.5 + 1.0 + 1.5 + 2.0 B-V

Figure 11. Hertzsprung Diagram for CPMLIST Binaries Having Published UBV Photometry. 41 Similarly Luyten's group B can be identified with pairs consisting

of a cool degenerate with (B-V) _> +0.5 and a late type dwarf of

(B-V) +1.3. Relatively few of these systems have WD components

bright enough to have accurate UBV observations; in fact relatively

few such cool degenerates have been found (cf. Liebert 1979).

Nearly all of the upper main sequence objects (H < +14) D have WD companions on the upper WD sequence (H„ < +18). This com- D — bination produces pairs showing magnitude and color differences

similar to group B pairs, hence are denoted as B' pairs in

Figure 11, although technically classified by Luyten as group

A pairs. It is the B' pairs having 6 <_ A m <_ 7 and -3 Ac j<-2

which account for the previously unexplained maxima in histograms

for Am and Ac (Figures 8a and 9a, respectively) and the P 8 distinguishability of this group appears to be a consequence of the

discontinuity in the main sequence near HD = +14, (B-V) = +1.3 D where the onset of TiO blanketing in the B bandpass is rapid.

The evidence for distinct types of WD binaries is most con­

vincing in the plot of AB vs. | A (B-V)| presented in Figure 12.

Here magnitude and color differences are as previously defined for

A m and Ac (compare this figure to Figure 10). Mean AB and PS | A(B-V)| values for Luyten group A and B objects have been computed from the H,,D and B-V values plotted in Figure 11. it is evident that the two groupings labelled A and B in Figure 12 correspond

well to Luyten's two groups. Group B pairs are found along the

steep, roughly linear sequence. Red dwarf pairs also lie along this

line; however all are characterized by A (B-V) > 0 while binaries 42

12

10

8 o.ERrAC

6

4 ©,ERI AB ERI+BC.

2

0

0.0 + 1.0 +2.0

Figure 12. Comparison of the AB and |a (B-V)| Values for All CPMLIST Binaries Having Published UBV Photometry, Showing the Empirical Separation of Luyten*s Group A and B, B' CPM Pairs. containing a late type dwarf and cool degenerate star (e.g. R193 +

VBsll) generally have A(B-V) < 0. Pairs of group A are less tightly

confined but seem to consist of systems like 0£ Eri BC. Values of

A(B-V) may be either positive or negative. Note the sparse region near the lower left corner of Figure 12 where double degenerate pairs would be located. The central region of this diagram contains

five systems which cannot be ascribed to either group. One object,

CoD-38°10980, is the hottest WD in the diagram, with the brightest primary of the sample. Another is a GB pair of uncertain physical association. The remaining three have K type primaries located at or near the apparent discontinuity in the main sequence (H„b - +14 in Figure 11). • It is interesting that the triple system O2 Eri can be classified three ways: o^ Eri BC (DA + dM6e) is representative of group A, O2 Eri AC (dKl + dM6e) is a red dwarf pair characteristic of the nondegenerate pairs contaminating the group B sequence, while

O2 Eri AB (dKl + DA) has an indeterminant status. Regression lines have been fit to the group B sequence which include (solid line) and exclude (dashed line) O2 Eri AC.

At this point it is worth noting that otherobservational selec­ tion effects favor the manifestation of two types of WD + RD binaries.

During the course of proper motion surveys it has been common prac­ tice to search the immediate field around a bright high proper motion object for low luminosity CPM companions. This introduces a bias towards the discovery of group B (or B') pairs (see Figure 11) since the bright primaries are generally G or K dwarfs and faint companions must necessarily be WD’s (or M dwarfs). Furthermore searches for 44

faint WD's have often emphasized faint CPM pairs of comparable

magnitude and strongly contrasting colors (e.g. the GB pairs). In

such cases a bias also exists for the detection of group A pairs

which normally consist of a WD and late MS companion. Thus the

distinct groups of WD + RD CPM binaries noted by Luyten now seem

ascribable to selection effects inherent in the search for WD's via

proper motion surveys and to the intrinsic properties of H-R diagrams

constructed from photometric and parallax data.

It is surprising that so many details of a conventional H-R

diagram remain discernible from the crude magnitude and color

estimates supplied by the proper motion catalogs. The real value of

the Luyten and Giclas surveys lies in the fact that such programs

have compiled rough photometric and colorimetric data for well over

100,000 stars in the solar neighborhood. CHAPTER III

INSTRUMENTATION AND REDUCTIONS

Silicon Intensified Target (SIT) vidicon spectra of 47 CPMLIST

pairs were obtained in June 1979 and February 1980 at the Cerro

Tololo Inter-American Observatory (CTIO) 1.5m telescope. Detailed

information on the operation and data reduction procedures for the SIT

vidicon is available in the CTIO Facilities Manual (Schaller et al.,

1978) and a general review of this instrumental system has recently

been published by Atwood et al., 1979). For these reasons only a brief

summary of procedures relevant to the reduction of the CTIO SIT vidicon

spectra is presented here.

The SIT vidicon is a two-dimensional multi-channel (MC) detector

which, in the spectroscopic mode, provides a 70 x 512 element image

with the long axis oriented along the direction qf the spectrograph

dispersion. Interactive reduction software exists at CTIO for handling

such array-oriented data (cf. Schaller et al., 1978). For example,

individual image lines parallel to the spectrograph dispersion may be

extracted and co-added whilq adjacent image lines of the night sky

spectrum may be subtracted from the accumulated stellar spectrum.

Individual spectra can be manipulated in a number of ways; one

capability is that overlapping spectra of close pairs of stars can be at least partially separated by careful weighting and subtraction

45 46

of spectral Image lines. Several close CPMLIST pairs required this

treatment. Flux and/or spectral line measurements for these close

pairs carried zero weight in the calibrations.

After the CPMLIST program spectra were extracted a wavelength

calibration was obtained from measurements of Fe-Ar and Hg emission

line spectra secured during each night of the observing run.

Corrections for atmospheric extinction were applied to all spectra

using tables of mean extinction coefficients for CTIO. Thus the

reduced spectra may contain uncorrected changes in extinction during

a given night. For each night the wavelength-calibrated spectra of

three or more standard stars from the lists of Oke and Schild (1970)

and Oke (1974) were used to determine an absolute flux calibration.

Because of cirrus, most of the ten nights of observation reported

here could not be considered suitable for absolute photometry.

This assessment was substantiated by residuals as large as + 10% in

the calibrated fluxes for individual standard stars. These residuals

are assumed to be representative of the spectrophotometric accuracy

of the program star spectra. In an attempt to minimize this source

of error relative fluxes or intensity ratios are preferentially used in the analyses which follow this chapter.

The photometric accuracy of the vidicon spectra is demonstrated by Figure 13 where an instrumental B ’ magnitude closely approxi­ mating the B magnitude of the UBV system (Johnson, 1952) has been computed from all well-exposed spectra using Equations III-2 and

III-3 given below and integrating the flux-calibrated spectra from

3800 - 5100 X. It is apparent from the plot that accuracy comparable 47

60

58 -

/

56 -

54

52

• WD 50 + RD

8 10 12 14 16 B

Figure 13* The B* vs. B Calibration for Degenerate and Non-Degenerate Stars Observed With the SIT Vidicon. to that obtainable by photographic means was obtained even on mediocre nights such as these. Note the similarity of the (B', B) relation for the WD's and late type MS companions having published

B magnitudes. A nightly determination of extinction under more photometric conditions would likely permit magnitude determinations of accuracy approaching that of conventional photoelectric photo­ metry while retaining full spectroscopic information. Comparisons to previous low-dispersion spectroscopy and/or filter photometry for an object can thus be made by convolving the appropriate resolution or transmission curve with the observed digital spectrum and inte­ grating over the desired bandpass.

Most CPM systems in the program consist of a faint WD plus red dwarf companion whose energy maxima lie at UV and near 1R wave­ lengths, respectively. Since measurements of blue spectral region features such as the Balmer lines and Hel lines in WD’s, and Call

H&K, Cal 4226, and TiO 4954 in the late type components were desired, a grating was chosen which provided a reciprocal dispersion of M.00 8/mm on the detector and a spectral resolution R a-9 8 (^3 pixels, FWHM) for spectra roughly spanning the B magnitude region

3700-5200 8. This configuration allowed spectra of objects as faint as B = +17 to be obtained in about 120 minutes.

The SIT vidicon and similar instruments have been extensively employed in studies of faint emission line objects such as planetary nebulae and active galactic nuclei. Relatively little has been published on the use of such devices for the determination of stellar absorption line strengths (cf. Greenstein and Vauclair, 1979; Hesser and Bell, 1980). One objective of this program will be to establish

the possible use of the SIT vidicon for quantitative absorption line

and colorimetric work on very faint stars. In the next section the

numerical techniques used to derive instrumental magnitudes, flux

ratios, colors, and equivalent widths from digital spectra are

discussed.

Reduction of MC Spectra

A raw SIT vidicon spectrum as recorded on magnetic tape

consists of a 512 element array containing the number of counts N

accumulated in each pixel number n. If the calibrations mentioned

above have been performed the spectra have units of observed flux

f (X ) per channel of central wavelength A . For a given wavelength n • n region the integrated flux can be defined in terms of the mono­

chromatic fluxes f^ summed over a finite number of channels

(n^ - iiy), each having width AAn :

A„ rL nR f (A )AA F - / fxdA = fv (An) [ ] dAn - c Z [ n L A. n=n,t A n=n„ A 1 V n V n III-l

If the dispersion curve is a slowly varying or constant function then AA^ = R 1, the average channel spacing, and Equation III-l may be written 50 In general R'< R, the actual instrumental resolution, since an in­

dividual pixel need not be an independent picture element. For

example R' = R/3 = 3 X for the SIT vidicon throughout the wavelength

region covered by the CPMLIST spectra.

Equation III-2 may also be used to define an instrumental magnitude

m = -2.5 log F . III-3

Similarly a flux ratio may be defined as

r12 ' Fl/F2 • FI1-4

Either of these two relations may be used to obtain an instrumental color index between two spectral regions

C12 = ml ” ra2 = ”2,5 log r12" III-5

In principle the determination of equivalent widths from digital spectra is also straightforward since

w ■ £ 1 1 ' O o IdA “ F { 1 1 - f (x ,o n l -6 A n=nv v n where fv(An>£.) and f^CA^jc) represent the observed monochromatic line and continuum fluxes, respectively. As before, AAr may be replaced by R' if the instrumental dispersion does not vary significantly over the integrated bandpass. Equation III-6 can be difficult to apply to very broad or very weak spectral features.

For example the wings of H6, He, and higher level Balmer lines overlap in DA WD's, rendering the placement of continuum points 51

arbitrary. The problem posed by weak spectral features will be discussed later. Additional difficulties imposed by the finite resolution and format of the instrumental output from modern MC detectors have been discussed by Gray (1976) and Greenstein (1979b).

The procedure adopted here for calculations of equivalent widths from the CTIO SIT vidicon spectra is similar to that used by Greenstein and Vauclair (1979). A continuum is numerically defined by a straight line in (v, log f^) between endpoints chosen to minimize contamination by overlapping line wings and/or nearby weak lines. Since the SIT vidicon spectra have units of (fy,A) Equation III-6 becomes

iR W = R' E [ 1 - dex (log f (A ,£) - log f (A±,c) ) ] . III-7 ±=1V

For the vidicon spectra an equally-spaced data set f^(A^,£) is interpolated from the observed spectrum >£) using the formula

A .-A. A f (A ,l) log t w o - log y i „ +1,i> + [ ^ log n+1 n i v n+1 III-8 which is linear in (v, log f ). This interpolation formula is also used for each corresponding continuum point, except that Ar = Ar and = A represent the endpoints of the desired integration region. Such an equally-spaced data set provides two distinct advantages over simple sums of the observed data points. First,

R' = constant and therefore can be removed from the sum performed in Equation III-7. Usually the number of interpolation steps is chosen so that the spacing between data points remains similar to that of the original spectra. Second, a consistent treatment of

weak spectral features which extend over only a few pixels is

possible since the interpolated spectral region can begin and end

precisely on the wavelengths specified. Endpoint channels in the

original spectral array generally straddle the desired endpoint wavelengths, sometimes contributing a significant amount of flux not contained in the desired wavelength region. Equivalent widths are not particularly sensitive to this difficulty as long as valid continuum endpoints have been chosen; however the problem can affect extracted narrow band flux sums, magnitudes, flux ratios, and colors.

A partial correction for this effect can be made if an equally-spaced data set is interpolated to match the desired endpoints.

The FORTRAN program VIDIC0N2 was developed to calculate the instrumental magnitudes, flux ratios, colors, and equivalent widths by the methods described above. A listing of this program is pre­ sented in the Appendix. It can be adapted to the analysis of MC spectral data sets other than those obtained with the CTIO SIT vidicon. CHAPTER IV

ABSORPTION LINE MEASUREMENTS IN THE SPECTRA OF DEGENERATE MEMBERS OF CPM BINARIES

Interest in the spectroscopic properties of the degenerate members of CPM binaries has seldom gone beyond the routine identification of WD spectra among the various lists of CPM pairs that consist of faint bluish stars having brighter yellow or red companions. The reasons for this lack of emphasis are two-fold.

First, though presumably coeval, the members of wide pairs are generally assumed to experience evolutionary changes which are not influenced by the existence of a companion; hence there exists no a priori reason for suspecting the basic properties of CPM WD's to be at variance with those of single WD's. The results of the present study substantiate this assumption. Second, in Chapter II it was shown that most CPM pairs contain at least one component which is considerably fainter than m = +15, making spectroscopy with con­ ventional equipment difficult. The large color differences often exhibited by the components of such pairs further impede most observing programs constrained by a single instrumental system. High quantum efficiency detectors such as the CTIO SIT vidicon now allow such programs to be undertaken. The purpose of .this chapter is to establish the extent to which this and similar instruments may be used for absorption line measurements in faint degenerate stars and

53 54 to establish an effective temperature (hereafter referred to as

"temperature" or = T^j/IOOOK) and gravity (log g) for each

CPM WD. These results are compared to previously published tem­

peratures and gravities for single and CPM WD's.

There exist two fundamental means by which the temperature and

gravity estimates for a WD may be determined. A two-color diagram

such as (U-B, B-V) or (u-b,b-y) can usually be used. In cases where

strong spectral features, e.g. the Balmer lines in DA WD's of inter­

mediate temperature (11 £ T^ ± 15), affect one or more photometric

bandpasses the derivation of temperature and gravity from the

observed colors may be difficult. However, photometry for even very

faint objects is relatively easy to obtain and it is therefore the

most common- information available for WD's beyond the proper motions

and color classes provided by Luyten or Giclas. The other, more

desirable means of determining the temperature and gravity of a WD

is to compare the observed spectrum to a suitably complete grid of model atmospheres. In principle this involves fitting computed

absorption line profiles to the observations. In practice, however, measurements of equivalent widths W and/or FWHM are used in con­ junction with photometric colors to derive the temperature and

gravity. It may be necessary to fit the entire spectral energy distribution with a model atmosphere for objects that exhibit featureless or peculiar spectra.

Ideally the temperature and gravity estimates derived from photometric and spectroscopic observations will agree. Unfortunately calibration uncertainties do not allow a direct comparison of 55 observed photometric colors to those predicted from model atmosphere calculations (cf. Shipman, 1979). Spectroscopic observations are therefore desirable, particularly when the data are from MC de­ tectors where digital techniques may be employed in the reductions and analysis. Only recently, however, have sufficiently complete grids of model atmospheres for the temperatures and gravities common to WD's been available in the literature (cf. Koester et al. 1979;

Koester, 1980). These models have been employed in the present study.

Table 3 presents general information on each CPM WD observed during this program. The first five columns contain respectively the earliest designation found in the literature, the Luyten or

Giclas catalog number from which an accurate 1950 position and/or finding chart was obtained (omitted if identical to the entry in column 1), the spectral type, and color class (B-V color is given when available). Photographic (or B) magnitudes, proper motions y, and reduced proper motions H (or H^) are also provided in columns

6-8. The photometric UBV data have been taken from U.S. Naval

Observatory publications (Harrington et al., 1978 and references therein) unless otherwise noted. Spectral types have been assigned using the criteria summarized by Greenstein (1960) and Hintzen

(1975). A few which appear to conflict with previous low-resolution classifications have been noted. Asterisks mark newly identified WD's.

The vidicon spectra have been analyzed by the numerical methods outlined in the last chapter. Since different absorption lines are measured in H-rich (mostly DA) and He-rich (mostly DB) WD's, the results for these objects are discussed separately. The objects Table 3. Degenerate CPM Stars Observed With the SIT Vidicon

Name Source a,9s0 6,950 Sp Color mpg>B y HpgjB Notes

*G272-B5B 02h00m31S -17°07l5 DAs -1 15.8 0.05 14.3 a 02 Eri B G160-60 04 13 00 -07 44.0 DA +0.03 9.55 4.08 17.60 b *LP891-12 04 43 18 -27 32.0 DB: f 15.6 0.24 17.5 c G102-39 05 51 05 +12 23.8 DC -0.12 15.80 0.28 18.15 BPM18164 L182-61 06 15 36 -59 11.4 DB -0.09 14.00 0.33 16.59 d *LP895-41 06 42 34 -28 30.8 DA a 16.8 0.16 17.8 L745-46A LP783-3 07 38 02 -17 17.4 DF +0.32 13.36 1.26 18.86 b LDS235B LP786-6 08 45 18 -18 48.0 DB -0.06 15.49 0.16 16.51 b *LDS275A LP462-56 09 35 00 -37 07.0 DC 1 [+0.13 14.43 0.37 17.27 e,f *LDS275B LP462-56 09 35 00 -37 07.0 DAsJ L753-52 LP791-55 10 43 30 -18 50.0 C2 g 16.6 1.98 23.1 g G163-B9B 10 43 39 -03 24.1 DAs +0.43 15.75 0.08 15.27 e,h L970-30 LP672-1 11 05 28 -04 52.8 DA +0.09 13.01 0.44 16.23 e *LP849-59 11 07 00 -25 43.0 DC: a 16.8 0.25 18.8 c *LP378-537 13 04 48 +22 43.0 DA f 16.2 0.11 16.4 W485 G14-58 13 27 40 -08 18.8 DA +0.09 12.40 1.24 17.87 LDS455A LP798-13 13 34 18 -16 04.0 DA -0.06 15.29 0.12 15.69 b *LP498-26 13 36 45 +12 23.8 DB b 13.9 0.19 15.3 L619-49 LP856-53 13 48 30 -27 19.0 DA a 15.1 0.24 17.0 *LP916-27 15 42 18 -27 30.0 DBwk f 15.5 0.24 17.4 -38°10980 Wg31 16 20 38 -39 04.7 DA -0.14 10.86 0.07 10.32 d *BPM24150 L266-196 16 23 00 -54 05.0 DA +0.10 15.84 0.08 15.36 i G154-B5B 17 43 04 -13 17.3 DA +0.30 14.59 0.09 14.36 e LDS678A L923-21 19 17 54 -07 45.0 DC: +0.07 12.35 0.20 13.82 b,j LTT15921 G24-9 20 11 32 +06 32.5 DC: +0.36 16.09 0.70 20.32 VBsll LHS3601 20 54 06 -05 03.0 DC +1.13 17.81 0.82 22.38 b LDS749B LP638-4 21 29 36 00 00.0 DB -0.13 14.60 0.41 17.72 BPM27606 L283-7 21 54 24 -51 14.0 C2 +0.16 14.84 0.40 17.85 e BPM44275 L427-60 21 54 48 -43 42.0 DB -0.13 14.91 0.22 16.25 i LDS785A L573-108 22 24 36 -34 27.0 DB a 13.9 0.21 15.5 *LP581-35 22 53 24 +05 30.0 DB: f 15.8 0.45 19.1 c LDS826A L577-71 23 51 31 -33 32.8 DAs a 14.5 0.50 18.0 G275-B17B 23 52 56 -25 33.0 DA 0 14.2 0.05 12.7

*: Denotes new WD spectroscopic identification, a: Gr520, recently announced by Greenstein (1980). b: UBV by Eggen and Greenstein (1965a, 1967). c: Very weak spectrum, d: UBV by Wegner (1973). e: UBV by Eggen (1968). f: Very close pair, UBV photometry of both components, g: Recently announced as WD by Liebert et al. (1979). h: Originally classified as "Gp" by Eggen (1968). i: UBV by Wickramasinghe and Bessell (1977). j : Originally classified as DAwk by Eggen and Greenstein (1965a). 57 which exhibit featureless (DC) or peculiar (C2) spectra require analysis of large spectral regions which could be affected by errors in the flux calibrations or by changes in sky transparency that occurred during most of the nights of observation reported here.

Shipman (1979) has computed model atmospheres for some of these objects and his temperatures, gravities, and compositions have been adopted here.

The H-Rich CPM WD's

Koester et al. (1979) have pointed out that the ultimate fate which awaits about 80% of all stars in the Galaxy is to become a DA

WD, characterized by strongly Stark-broadened hydrogen line spectra.

Despite a large range in temperature (7 < < 70) and luminosity

(-4 < log L/L0< 0) Koester et al. (1979) and Shipman (1979) have shown that DA WD's are a surprisingly homogeneous group in terms of mass and chemical composition.

As indicated by the previous discussion at least two obser­ vational parameters are required to uniquely specify the temperature and gravity of a DA WD from MC observations. In principle the star's position in a properly calibrated two-color diagram will suffice

(Greenstein, 1976). In practice, however, accuracy is improved if one color is chosen to be particularly sensitive to temperature and the other is replaced by the measured equivalent width of an absorption line such as Hy, which is fairly sensitive to gravity for Tg _> 13. MC detectors are especially suitable for this purpose since the resulting output can provide both color indices and equivalent widths (cf. Greenstein, 1979a, b). In addition the measurement of absorption line strengths and widths is much more

straightforward than for photographic spectra. Such measurements

are not likely to be seriously affected by the sources of error which influence broad baseline MC colors. In addition, since photoelectric colors are not available for all the DA objects in the present sample, a means of evaluating temperatures and gravities from the spectra alone would be desirable. Two approaches to this problem will be examined which utilize the sensitivity of the Balmer lines to both temperature and gravity in DA atmospheres. Schulz and Wegner (1981) provide a review of the mechanisms affecting hydro­ gen line profiles under the conditions found in WD's. The observa­ tional manifestations of these effects are briefly summarized below.

The absorption lines of HB and Hy reach maximum strength in WD spectra in the temperature range 11 <_ Tg <_ 15 for 7 log g 9 respectively. At cooler temperatures the line strengths weaken but become relatively insensitive to gravity. Thus gravities may not be determinable for cool WD's from line strength measurements alone. In this regime MC spectral indices such as (u-b) and (g-r) may be used to estimate log g (Greenstein, 1976; Shipman and Sass, 1980). Above

Tg = 15 the lines also weaken. Here temperatures and gravities may both be determined from line strengths alone (the decision as to whether a particular WD belongs to the cool regime (Tg < 13) or the hot regime (Tg > 13) can usually be made by careful inspection of the line profiles (cf. Weidemann and Koester, 1980) or by the colors computed from MC spectra (cf. Greenstein, 1976)). In the above discussion only total line absorption, character­ ized by the equivalent width W, has been considered. Another useful quantity is the width of a Balmer line, conventionally expressed as full width at half-maximum intensity FWHM, which usually varies in the same general way as W under different temperature and gravity conditions. However the computed FWHM for the atmospheric models provided by Koester et al. (1979) reveal a much stronger gravity dependence than that exhibited by W in the hot temperature regime.

Because of this effect measurements of W and FWHM provide a means of deriving temperatures for cool DA's, and both temperature and gravity estimates for hot DA WD’s. These estimates appear to be of comparable accuracy to more conventional techniques which require both spectro­ scopic and photometric observations.

Figure 14 illustrates this method for Hy. Computed values of

W(y) and FWHM(y) for each model atmosphere published by Koester et al.

(1979) have been plotted for the hot (a) and cool (b) DA temperature regimes, respectively. Lines of constant Tg (dashed) and log g (solid) are labelled. Also plotted in this figure are the individual FWHM(y) and W (y) values measured for the DA objects from the vidicon spectra

(small dots indicate low quality spectra). The above-mentioned effects are readily apparent; temperatures and gravities are obtainable for hot DA WD's, though gravity effects are indiscernable in cool objects.

Schulz and Wegner (1981) have shown that higher order Balmer lines such as H6 and He, whose wings strongly overlap, weaken with increasing gravity in contradistinction to the previously discussed behavior of HB and Hy. This result is readily understood as a Figure 14. The Dependence of The Dependence FWHM(y) and on W(y) Temperatureand Figure14.

FWHM (T) FWHM (T) 20 40 30 20 50 30 60 Gravity in DA inGravity DA WD's. F110, • High Wt. • Lew Wt. 0 0 0 40 30 20 10 50 40 30 25 25 20 30 50 40 W(T) G154-B5B \11 13 61 consequence of the fact that no true continuum exists between over­

lapping lines, hence reference to pseudo-continuum points in the

measurement of equivalent widths will naturally yield smaller values

for higher gravity atmospheres where Stark broadening is more severe.

This behavior provides a second means by which temperatures and

gravities of DA WD's may be derived. Figure 15 displays the

relations between the mean of the equivalent widths W(B) and W(y),

W (66), and that for W(6) and W(e), W(6e), computed by Schulz and

Wegner (1981) for the grid of hot (a) and cool (b) DA atmospheres

provided by Koester et al. (1979). As in Figure 14, lines of constant

temperature and gravity are indicated, and loci for the DA WD's

observed with the vidicon are marked.

The W's and FWHM measured from the SIT vidicon spectra for

H-rich CPM WD's are presented in Table 4. Where multiple obser­ vations of an object were obtained an internal standard error is

also listed. When, as here, such measurements are to be compared

to theoretical values the instrumental resolution R, exposure time,

and choice of continuum points must be considered.

Greenstein (1980) has summarized the effects of finite in­

strumental resolution on the theoretical line profiles of hydrogen-

line spectra. The sharp cores present in most DA spectral lines are broadened to widths comparable to the instrumental resolution, producing a larger residual intensity at the line core than the

theoretical value. Hence the observed FWHM for such lines, here­ after designated as FWHM', will be appreciably larger than that predicted by model atmospheres. On the other hand the broad wings 62

<*> T, > 13

*5 . W(se) li

j o 10 2%

JO • High Wt, FflO- 40, "40 100CW •0 • Low Wt. 100 7S I

20

W(se)

10

10 20 30 40 W(0T)

Figure 15. The Dependence of W(6e) and W(8y) on Temperature and Gravity in DA WD's. Table 4. Absorption Line Measurements in the Spectra of DA White Dwarfs.

W'e W'6 W'Se FWHMy a Name n ae W'y o W'B a : W ’8y cfc T3 logg T3 logg T3 logg +50 A +50A +100A Y +100A p +50A +100A PT +100A Y Vidicon Koester Shipman

G272-B5B 1 15.9 11.6 18.8 14.3 13.8 16.6 13.9 8.5 7.0: 02 Eri B 8 15.2 1.3 19.8 0.9 34.3 0.9 27.4 1.8 17.5 0.5 30.8 1.0 39.8 2.3 18.2 7.7 16.9 7.6 16.9 8.0 LP895-41 1 10.6: - 13.2: - 26.5: - 13.4: - 11. ft - 20.0: - 25.0: - 9.3: 8.3: -- - - LDS275B 1 4.4: - 7.4: - 11. a - 4.9: - 5.9: - 8.a - 12.6: - 7.8: 8.a 8.2 9.4 - - G163-B9B 2 7.1 0.7 3.5: 0.4 7.8 1.2 3.8 3.3 5.3: 0.1 5.8 2.3 4.8: 2.7 7.2: 7.0: - -- - L970-30 10 16.4 1.7 20.6 1.1 36.1 1.3 30.7 1.3 18.5 1.3 33.4 0.8 44.3 2.8 17.8 7.9 15.4 8.0 15.3 8.0 LP378-537 2 5.2: 0.1 21.4 0.2 35.4: 1.5 26.0c 2.8 13.3 0.1 30.7: 2.1 29.6 0.4 10.4: 9.0: --- - W485 9 15.8 0.9 20.1 0.9 39.3 1.0 35.0 1.7 18.0 0.6 37.2 1.1 47.7 1.8 16.8 8.2 14.0 8.0 12.8 7.7 LDS455A 2 12.0 0.2 17.1 1.0 33.4 0.1 29.1 0.7 14.6 0.6 31.2 0.4 44.2: 0.4 20.6 8.3 21.0 8.0 20.3 8.6 L619-49 1 13.1 - 14.2 - 31.6 - 21.2 - 13.6 - 26.4 - 34.a - 10.a 8.6: - --- -38°10980 2 8.5 0.3 14.0 0.4 26.5 0.8 21.3 1.1 11.2 0.0 23.9 0.9 35.8 0.5 26.5 8.1 24.5 7.4 -- BPM24150 1 13.0 - 18.6 - 36.0 - 29.7 - 15.8 - 32.8 - 40.9 - 9.4 8.6 -- -- G154-B5B 1 14.3 - 18.4 - 34.4 - 34.0 - 16.4 - 34.2 - 43.0: - 9.9 8.5 10.0 7.5 10.6 7.2 *F110 2 3.8 1.2 2.6c 0.4 7.2 1.8 -1.4: 0.5 3.2: 0.4 2.8: 0.6 14.8: 1.6 50: <7 38.8 6.0 -- LDS826A 1 8.2 - 9.7 - 25.9 - 14.5 - 9.0 - 20.2 - 27.7 - 9.4 9.2: - - -- G275-B17B 1 9.9 - 9.4 - 15.4 - 11.6 - 9.6 - 13.5 - 15.9 - 8.4 7.7 - -- -

*: sdO used as SIT vidicon flux standard. Observed by Greenstein and Sargent (1974), from which the values of T^ and log g listed in the next to last column were taken. Estimates for vidicon T^ and log g are extrapolated from Koester's grid of model atmospheres.

LOO' 64 of these lines are relatively unaffected. W's are therefore seldom seriously affected by instrumental resolution as long as true con­ tinuum points are chosen.

Figure 16 exhibits the correlation between the theoretical and observed Hy equivalent widths (W(y) and W'gj^y), respec­ tively). Values of W(y) were interpolated from tables provided by

Koester et al. (1979) for several objects having prior temperature and gravity determinations. To obtain better coverage at small equivalent widths several points representing vidicon measurements of the Hel 4026 and Hel 4472 lines in DB stars have also been plotted.

For these objects theoretical W ’s have been interpolated from tables published by Koester (1980). A line of unit slope is indicated in the plot and it is apparent that agreement between observed and theoretical equivalent widths is good to Md.0%. In contrast the agreement between Wgj^(y), equivalent widths determined with the

CTIO SIT vidicon, and W ’ (y), those derived from photographic PS spectra, is noticeably worse. This comparison is made in Figure

17 where W g IT(y) is plotted against W' (y) from the lists of

Eggen and Greenstein (1965 et seq.). Again some equivalent widths of Hel lines in DB WD's have been added to fill out the lower end of the relation. Several points fall well off the unit slope relation in the sense that W' (y) exceeds W' (y). This tendency was PS also noted by Greenstein and Vauclair (1979) when examining the re­ lation between photographically-determined equivalent widths and those determined from Palomar SIT vidicon or MC spectra.

It is tempting to speculate that the larger errors exist among the W ’ . One likely source of difficulty is underexposure which Figure Comparison16. theof Instrumental and Theoretical H,He Equivalent

V) (H,He) (A) 40 20 30 Widths forWidths and DA DB WD'sObservedthe SIT With Vidicon. 10 20 30 •#o •#o 40 • DA • O DBO LowWt.

65 66

40

30

20

• DA O DB

20 30 40 Wpg (T)

Figure 17. Comparison of SIT Vidicon and Photographically Determined H,He Equivalent Widths for DA and DB WD's. causes the cores of strong Balmer lines in faint WD’s to be lost in plate noise, resulting in underestimated equivalent widths. Vidicon spectra need not be subject to this problem since multiple exposures may be co-added. However equivalent widths will be underestimated in saturated vidicon exposures since the effect is to depress the con­ tinuum level. Fortunately when this situation exists it is evident as an overflow in one or more channels before sky levels are -sub­ tracted. Wegner and Schulz (1981) found that many DA WD's previously thought to exhibit abnormally weak Balmer lines as determined from earlier unintensified photographic spectra in fact have relatively normal line strengths when reobserved with more sensitive detectors which are now available. Such findings continue to support the re­ sults of Koester et al. (1979) which imply a rather narrow mass dispersion among DA WD's (0.58 + 0.10 M ).

As mentioned earlier, measurements of equivalent widths in DA

WD's are particularly sensitive to the observer's choice of con­ tinuum points since in many cases the line wings overlap.

Weidemann and Koester (1980) and Koester (1980) stress the importance of using continuum points which are consistent with those used to calculate theoretical equivalent widths from model atmospheres.

The pseudo-continuum points used for H and He lines in these two papers have been adopted here in the measurement of the vidicon spectra and are listed in Table 4. As one example, the effects of an alternate choice of continuum points for Hy are illustrated in Figure 16. Arrows indicate the resulting changes in W' T (y) bXl when more widely separated continuum points (4220 _< £ 4500 8) 68

used by Greenstein and Vauclair (1979) are employed. The differences

in Wgjj(y) < 30 X are too small to show presumably since little wing

overlap occurs in Hy for objects having relatively weak lines. How­

ever wgjT (Y) f°r strong-line W D ’s are up to 10% larger, as expected.

The fact that these values also appear to improve the correlation

between the observed and theoretical equivalent widths is probably

fortuitous considering the small number of points involved.

Although no significant corrections to the vidicon equivalent widths appear to be necessary such is not the case for the observed

line widths, FWHM1. The fact that the observations do not fully resolve the sharp line cores present in the theoretical line profiles must be taken into account. Ideally the instrumental pro­ file should be convolved with each theoretical line profile in the grid of model atmospheres spanning the temperature and gravity range of WD's. From these a set of predicted FWHM' can be obtained which may be directly compared to the observed FWHM'.

Unfortunately this procedure requires line profiles for each model atmosphere which are not generally available. We adopt here an empirical approach which appears to provide a satisfactory scheme by which corrected FWHM can be obtained.

Weidemann and Koester (1980) have recently published, in graphical form, theoretical Hy line profiles for the DA WD atmo­ spheres originally computed by Koester et al. (1979). These plots provide insufficient numerical detail for convolution of in­ strumental profiles; however it is possible to measure FWHM to an accuracy of about + 1 X by eye. Previously mentioned was the fact that observed line cores are shallower than those predicted by

theoretical profiles and appear broadened to widths comparable to

the instrumental resolution. As a crude approximation to these

effects a set of predicted FWHM' was obtained by measuring each

theoretical line profile computed by Weidemann and Koester but

with maximum absorption defined by an artificial flat line core

of residual intensity equal to that at the full width = 9 X level.

This procedure ignores the sharp line core in a manner similar to

that caused by the 9 X resolution of the vidicon spectra. Pre­

dicted FWHM' were obtained for all temperatures and gravities in

the grid of models below = 40. At hotter temperatures Weidemann

and Koester1s plotted line cores overlap and proved too difficult

to distinguish individually. Fortunately all W D ^ in the present

sample are cooler than this limit. This procedure neglects the small amount of line core absorption which in principle appears as a slight additional broadening of the line core but as will be seen the error introduced by this omission appears to be minor.

The above measurements were used to derive the linear least- squares correction formula

FWHM = 0.94 FWHM1 - 9.6 IV-1 which is illustrated in Figure 18. The 31 FWHM1 measured from the plotted line profiles exhibited a standard error of + 1.2 X from Equation IV-1, nearly all of which can be ascribed to the measuring accuracy mentioned earlier. Plotted in Figure 18 are points corresponding to several DA WD’s observed with the vidicon which have computed atmospheres by Koester et al. (1979) that were Figure 18. The Effect of Finite Instrumental Resolution on SIT Vidicon Vidicon SIT on Resolution Instrumental Finite of Effect The 18. Figure FWHM's|T(r) (A) 20 30 40 60 50 WM() nD ’s. D W DA in FWHM1(y) F110 10 FWHM(r) FWHM(r) 20 30 (A) 40 70 used to interpolate theoretical FWHM. Two of these stars appear to

have observed FWHM' which are appreciably larger than that pre­

dicted by Equation IV-1. Both have previously published FWHM,

indicated by arrows, that fall closer to the expected line. The

position of F110, a well-known sdO which falls outside the gravity

range considered here, is included for comparison. Also indicated

in Figure 18 is a regression line based on the five remaining

objects having previously published FWHM that agree with those inter­

polated from the models. Although coverage is poor at small FWHM

this relation appears to be in rough agreement with Equation IV-1.

An additional check on the empirical correction scheme re­

presented by Equation IV-1 was performed using the results of

Greenstein's (1980) examination of the effects of instrumental

resolution on theoretical Hy line profiles for several model atmo­

spheres between 15 <_ T^ ± 60 and 4 £ log g <_ 8. Greenstein con­ volved Gaussian "instrumental" profiles of FWHM = 4, 8, and 16 X

with each theoretical line profile. The resulting line profiles

were plotted at a scale sufficient to obtain measurements of FWHM'

and central residual intensity Plots of FWHM' vs. instrumental

resolution were used to interpolate the difference between the

"observed" FWHM' at 9 X resolution characteristic of the vidicon

spectra and the theoretical FWHM for each model. The mean correction factor derived for these models was = 6.9 + 3.3 X which is within lo of the average correction required by Equation IV-1. Two 72 WD line profiles considered by Greenstein (T^ = 60, log g = 8 and

Tg = 40, log g = 7) lack the sharp cores exhibited by the two

corresponding line profiles computed by Weidemann and Koester. The

latter are 6% and 9% deeper at the core, respectively. It is in­

teresting that consideration of only the profiles exhibiting sharp

line cores in Greenstein's sample results in = 9.2 + 1.4 X.

This value is much closer to the average correction required by

Equation IV-1. Thus it appears that the observed vidicon FWHM' can

be suitably corrected for the effects of instrumental resolution.

Equation IV-1 has been applied to all raw vidicon FWHM' to obtain

the FWHM listed in Table 4.

The W' and FWHM listed in Table 4 do not uniquely determine

the locus of a particular WD in Figures 14 and 15. The two methods of deriving temperatures and gravities illustrated by these figures both require an initial decision concerning the placement of an object in the "hot" or "cool" regimes indicated. To address this problem instrumental magnitudes m(40), derived from the summed flux between 3700 - 4200 X according to Equation III-3, and m(49), similarly defined between 4600 - 5200 X, were computed from the vidicon spectra and used to define a crude color index C(40-49).

Despite the photometric uncertainties discussed in Chapter 3 and the fact that the wavelength region spanned by C(40-49) is con­ siderably smaller than that spanned by U-V, the two color indices are fairly well-correlated (see Figure 19). C(40-49) can there­ fore be used to predict an approximate U-V for the DA WD's lacking

UBV photometry which can be used to decide whether an object belongs Figure 19. The Correlation Between SIT Vidicon C (40-49) and and (40-49) C Vidicon SIT Between Correlation The 19. Figure C C (40-49) + +0.4 0.2 0.4 0.0 0.2 U-V Color for WD's. for Color U-V F110 LP63B-4 1.2 0.8 U-V L283-7 0.4 0.0 s>y O DB O • DA • O DC O G24-9 73 74

to the "hot" or "cool" regimes using the (T^, U-V) calibration

provided by Koester et al. (1979).

The determination of temperatures and gravities thus becomes

a two-dimensional interpolation problem between the (FWHM(y),

W(y) ) plane, or (W(3y), W(Se) ) plane, to the (T^, log g) plane.

The curvature evident in the lines of constant T^ and log g in

Figures 14 and 15 make a non-linear interpolation scheme very

desirable. Collins (1981) kindly provided a general algorithm for parabolic interpolation in two dimensions from which a FORTRAN pro­

gram was written to interpolate the (T^, log g) values listed in

Table 4. The temperatures and gravities listed for each DA WD

in this table are the mean values predicted by each of the two methods discussed earlier in this chapter. Koester et al. (1979) and Shipman (1979) have independently determined (T^, log g) esti­ mates for several of these objects and these values are listed in

Table 4 for comparison.

The He-Rich CPM WD's

The procedures employed in the determination of temperatures and gravities for DB WD's, which exhibit absorption lines of Hel, are essentially those used in the analysis of DA spectra. In principle photometric colors can be used to derive (T^, log g) estimates. Unfortunately most DB stars are considerably hotter than T^ = 12. In many cases color indices formed in the visual wavelength region do not provide sufficient leverage for the determination of accurate temperatures and gravities. For example DB models computed by Koester (1979) for 12 <_ T^ 30 exhibit UBV

colors in the ranges -1.1 U-B <_ -0.9 and -0.1 £ B-V <_ 0.0. Con­

sequently only rough estimates of (T^, log g) can be obtained from

such data. It is therefore desirable to use the line spectra to

derive this information. Until recently, however, sufficiently

complete grids of model atmospheres and corresponding Hel line

strengths covering the range of temperatures and gravities common

to DB WD's were not widely available. Koester (1980) has published

such theoretical tables. These have been used in the analysis which

follows for the DB spectra obtained with the vidicon. The con­

tinuum points used for these theoretical W's have also been adopted

here.

Since the standard errors characteristic of the equivalent width

measurements of objects for which more than one vidicon spectrum was

obtained appear to be on the order of a = + 1-2 X the line(s) chosen

for the determination of temperatures and gravities should in most

cases exceed this strength. The strongest two Hel lines in the

vidicon spectra which meet this criterion are Hel 4026 and Hel 4472

whose functional dependences on temperature and gravity are illus­

trated in Figure 20. Unfortunately it appears that differences in

gravity are only marginally detectable from the vidicon spectra.

The equivalent widths of Hel lines in these objects were measured by the procedures discussed in Chapter 3 and the results

are listed in Table 5. A comparison of the equivalent widths of

Hel 4026 and Hel 4472 for objects in common with Wegner (1981) has

already been included with the results for DA WD's (see Figure 16) 76

logg 25 7.0. 7.5 4 8.0 H e l 4472

20

logg 7.0« 7.5 «

8.0 * H e l 4026

32 28 24 20 16 12 T3 (Koester, ‘80) Figure 20. The Dependence of W(HeI 4026) and W(HeI 4472) on Temperature and Gravity in DB W D ’s. Table 5. Absorption Line Measurements in the Spectra of DB White Dwarfs.

W' W' W' W' W f W' Name n 3gg9 0 4Q26 0 43gg 0 44?2 0 4713 0 4922 0 T3 log g

LP891-12 1 1.9: ------4.6: - - - <12: 8.0 BPM18164 4 1.6 0.1 6.1 0.8 - 11.5 0.9 2.6 0.6 6.9 1.4 14.8 7.6 LDS235B 3 2.8 0.3 7.1 0.6 - - 21.0 2.9 2.7 0.4 6.4 1.5 17.1: 8.4 LP498-26 1 3.1 - 7.5 - 1.6 - 12.7 - 1.9 - - - 15.0 7.0 LP916-27 1 0.5: - 1.7: - 0.9: - - - 0.4: - - - 12.3: 8.0 LDS749B 1 5.6 - 5.9 - 0.4 - 5.2 --- 3.6: - 13.8: <7: BPM44275 1 5.4 - 4.4 - 2.3 - 14.1 - 3.0 - 6.6 - 15.4: 8.5 LDS785A 1 3.6 - 6.1 - 1.9 - 18.2 - 2.0 « _ 7.4 - 16.1 8.3 LP581-35 1 - - 1.9: - 5.4: - -- -- 19.7: - 12.5: 8.0 and it appears no corrections to the vidicon instrumental values are necessary. As was previously done for the H-rich stars, Table

5 lists an internal standard error (a) when more than one spectrum of an object was obtained. Several of the faintest stars in this table have been noted as uncertain DB WD's because the noise present in such weak spectra admit the possibility that some are in fact

DC or DA WD's with very weak lines. The blue continua exhibited by these objects and the apparent absorption features at some of the wavelengths of Hel lines make the DB classification listed the best preliminary choice.

Koester (1979) found =* 7.4 for 25 DB WD's analyzed by the strengths of Hel 4026 and Hel 4472. From a larger sample of stars

Shipman (1979) found no evidence that He-rich WD gravities were system­ atically lower than those of H-rich WD's (for which = 8.0).

With this caveat log g = 8.0 (corresponding to one of the curves in

Figure 20) has been assumed and temperatures have been derived for the DB WD's in Table 5 having SIT vidicon equivalent widths which are insufficient for a gravity determination.

It should be noted that regions where W(4026) or W(4472) is increasing rapidly with temperature in Figure 20 are those positions where the gravity sensitivity of these lines is strongest. Apparently most of the objects listed in Table 5 have temperatures which fall in this region (12 _< T^ ^ 18). For the objects having reliable measurements of both Hel 4026 and Hel 4472 a theoretical (W(4026) vs.

W(4472)) plot was constructed following the procedure devised by Koester 79

(1979) from which the two-dimensional interpolation scheme described

in the last section could be used to derive (T^, log g) estimates.

Conclusion

It appears that fairly accurate equivalent widths and FWHM can be measured from the vidicon spectra. In addition temperatures and gravities for DA and DB objects are in most cases obtainable from absorption line measurements alone. Such absorption line measure­ ments with the vidicon are extended to the non-degenerate companions of these WD's in the next chapter. CHAPTER V

THE SPECTRA OF NON-DEGENERATE MEMBERS OF CPM BINARIES

One of the most vexing obstacles to a comprehensive understanding

of CPM binaries which contain a WD component is the scarcity of spec­

tral data available for the non-degenerate companions. This problem

is largely a result of the special emphasis which has historically

been placed on the degenerate objects. Over half the members of the

47 CPMLIST pairs observed with the SIT vidicon are late type red dwarfs

(RD's). These spectra provide a statistically significant data base

from which some general spectroscopic properties of such objects may be determined. In this chapter we discuss the results of a quantita­

tive study of the strengths of Call H&K, Cal 4227, H8, and the TiO

4954 (0,0) bandhead which are the most prominent spectral features exhibited by late type dwarfs in the 3700 - 5200 X wavelength region spanned by the vidicon spectra.

Table 6 contains the basic catalog data available for the CPM

RD's. The description of each column in this table is identical to that of Table 3 for the WD components. In Table 6 some B-V colors for CPM RD's lacking published UBV photometry and all spectral types listed have been estimated from the spectrophotometric indices derived from the vidicon spectra. Estimated B-V colors are annotated by a semi-colon.

80 81

Table 6. Non-Degenerate CPM Stars Observed With the SIT Vidicon.

Name Source a,950 6,95<) Sp B-V mpg>B p Hpg>B Notes

G272-B5A 02 ^OO^l8-!? ’07.5 dF8 +0.53: 12.5 0.05 11.0 -31°1454 L516-8 03 32 36 -31 14.0 dK2 +0.87: 11.8 0.51 15.3 LP888-25 03 32 19 -31 14.4 dM3 +1.47: 15.9 0.51 19.4 a 02 Eri A HD26965 04 13 00 -07 44.0 dKl +0.81 5.24 4.08 13.7 b 02 Eri C LTT1909 04 13 05 -07 44.4 dM6 +1.66 12.83 4.08 20.88 b,c LP891-13 04 43 21 -27 32.0 dM2 +1.35: 15.9 0.24 17.8 a +12°937 HD5998 05 51 05 +12 23.8 dF8 +0.58 8.34 0.28 10.58 -59°1275 HD44120 06 15 36 -59 11.4 dGO +0.58 7.05 0.33 9.60 d -28°3361 06 42 34 -28 30.8 dK3 +0.94: 11.2 0.16 12.2 -18°2482 LDS235A 08 45 18 -18 48.0 dK4 +1.00 12.63 0.16 13.65 b G114-B8A 08 58 17 -04 10.4 dF7 +0.58 10.86 0.10 10.86 G114-B8B 08 58 19 -04 11.2 dK4 +0.94 15.46 0.10 15.46 e -31°7352 HD82342 09 28 20 -31 53.2 dK4 +0.98 9.36 0.34 12.02 f LP902-30 09 28 20 -31 53.4 dM5 +1.46 14.54 0.34 17.20 f L753-40 LP790-30 10 37 54 -19 07.0 dM3 +1.48: 13.8 0.70 18.0 L753-39 LP790-31 10 37 54 -19 07.3 dM5 +1.60: 15.0 0.70 19.2 -18°3019 L753-52 10 43 30 -18 50.0 dM3 +1.40 12.42 1.98 18.90 g G163-B9A 10 43 39 -03 24.1 dF6 +0.62 12.03 0.08 11.55 h L970-27 LP672-2 11 05 34 -04 57.2 dM5 +1.52 14.07 0.44 17.29 h -25°8487 11 07 00 -25 43.0 dGl +0.60: 9.3 0.25 11.3 i L1046-18AB LP554-64/3 12 14 18 +03 14.0 dM4 +1.62 14.90 0.70 19.13 d,j +23°2539 13 04 48 +22 43.0 dG8 +0.81: 9.8 0.11 10.0 W485B G14-57 13 27 29 -08 26.7 dM6 +1.66 15.84 1.24 21.31 LP498-25 13 36 40 +12 24.7 dM4 +1.50: 14.5 0.19 15.8 L619-51 LP856-54 13 48 30 -27 19.0 dM3 +1.47: 13.9 0.24 15.8 LP74.0-11 14 23 30 -15 31.0 dF7 +0.51: 13.6 0.11 13.8 LP740-12 14 23 46 -15 22.6 dF4 +0.41: 15.5 0.11 15.7 LP916-26 15 42 16 -27 29.2 dM5 +1.58: 16.3 0.24 18.2 LP684-1 15 54 00 -04 41.0 dMO +1.20: 12.7 0.32 15.2 LP684-2 15 54 00 -04 41.1 dM4 +1.50: 15.5 0.32 18.0 -38°10983 HD147513 16 20 38 -39 04.7 dGl +0.63 6.01 0.08 5.47 k L266-195 16 23 00 -54 05.0 dMl +1.29: 14.1 0.08 13.6 -8°4352AB GL644AB 16 52 48 -08 14.7 dM4 +1.60 10.58 1.18 15.94 d,h -8°4352C GL643 16 52 45 -08 13.9 dM6 +1.70 13.40 1.19 18.76 h G154-B5A 17 43 04 -13 17.3 dM2 +1.44 13.35 0.09 13.12 h +9°3501 G140-B1A 17 50 33 +09 49.0 dK2 +0.95 10.31 0.10 10.31 h -4°4636A G22-9 18 53 42 -04 28.6 dGl +0.70 10.80 0.29 13.11 h,l -4°4636B G22-8 18 53 42 -04 29.1 dG7 +0.95 14.61 0.29 16.92 1 LDS678B L923-22 19 17 53 -07 44.7 dM6 +1.63 13.75 0.20 15.26 L1142-88 G24-10 20 11 32 +06 32.5 dMl +1.54 14.73 0.70 18.96 e,m R193 20 54 06 -05 03.0 dM6 +1.49 13.36 0.82 17.93 b,n -0°4234 LDS749A 21 29 36 00 00.0 dK3 +0.96 10.85 0.41 13.91 1 -51°13128 LDS765A 21 54 24 -51 14.0 dM5 +1.51 11.87 0.40 14.88 h,o LDS766B L427-61 21 54 48 -43 41.9 dM5 +1.54 16.08 0.22 17.79 P LDS785B L573-109 22 24 36 -34 26.9 dM2 +1.41: 14.1 0.21 15.7 82

Table 6. (Continued) O r t Name Source Sp B-V P 2J “l950 ^1950 mpg,B Hpg»B

LP581-36 22h53m24s+05°30.,0 dM2 +1.45 12.71 0.45 15.98 q G275-B8A 23 05 45 -29 04.7 dF9 +0.53: 16.3 - - r G275-B8B 23 05 45 -29 04.5 dK3 +0.96: 16.8 -- r LP933-66 23 09 34 -27 37.6 dM6 +1.64: 16.4 0.21 18.0 a,e LP935-14 23 44 36 -26 39.9 dMl +1.30: 11.7 0.36 14.5 G273-B1A 23 50 54 -08 21.1 dG8 +0.63 11.91 0.12 12.31 s LDS826B L577-72 23 51 31 -33 32.7 dM4 +1.55: 15.0 0.50 18.5 G275-B17A 23 52 56 -25 33.0 dF9 +0.54: 13.2 0.05 11.7 a: Very weak spectrum. b: UBV by Eggen and Greenstein (1965a,b, or 1967). c: Listed as dM5e by Eggen and Greenstein(1965a). d: Very close pair, spectrum likely contaminated by companion. e: Faint optical companion. f: UBV by Gliese and Jahreiss (1979). g: UBV by Eggen (1979a). h: UBV by Eggen (1968). i: Listed as dG5 by Luyten (1974). j: Uncertain identification, companion not seen. k: UBV by Wegner (1973). 1: Listed as sdK, idK by Eggen and Greenstein (1965b). m: Listed as dM5 by Eggen and Greenstein (1965a). n: Listed as dM4e by Eggen and Greenstein (1967). o: Listed as dMO by Luyten (1957). p: UBV by Wickramasinghe and Bessell (1977). q: UBV from Giclas et al. (1971). r: No detectable Lowell proper motion. s: Listed as dG3 in Harrington et al. (1978). 83

The bandpasses over which monochromatic fluxes were integrated

to form a set of line and pseudo-continuum indices for Call H&K, Cal

4227, HP , and the TiO 4954 (0,0) bandhead are summarized in Table 7. 2 Exposure times were generally chosen to yield at least 10 counts in

each 9 X resolution element ('v 3 pixels) at Call H&K depending on the

brightness of the particular star. Count rates near the red end of 3 these spectra usually exceeded 10 per resolution element. Photon

noise in the computed spectral indices is thus expected to be less

than 10% in all but the faintest stars observed. Table 8 presents

a summary of the computed spectrophotometric indices for the features

measured in the CPM RD spectra. Where more than one observation of

a star was obtained internal standard errors (a ) are recorded.

An additional 26 bright (m pg < 10) F, G, K, and M dwarfs were routinely observed throughout the program to serve as standards for

calibrating the observed spectral indices as functions of spectral

type and photometric color. These observations are recorded in Table

9. The final column in this table contains 1980 estimates of the

strengths of Call H&K emission (HKe) derived from long term observa­

tions reported by Wilson (1978) for objects in Table 9 which have

been shown to exhibit constant or quasi-periodic changes in HKe on

timescales of several months or years.

Quantitative Classification of Late Type Stellar Spectra: The Cal 4227 and TiO 4954 Indices

In stars of spectral type later than G5 the strength of the

Cal 4227 absorption line becomes very sensitive to temperature, reaching maximum strength in late M stars (Keenan and McNeil, 1976). 84

Table 7. Wavelength List for Features Measured in the Spectra of Non-Degenerate CPM Stars

Central Full Wavelength Width Feature (X) (X)

3901.1 20 Violet Pseudo-Continuum for Call H&K

3933.7 9 Call K2

3968.5 9 Call H2

4001.1 20 Red Pseudo-Continuum for Call H&K

4167.0 30 Violet Pseudo-Continuum for Cal 4227

4226.7 30 Cal 4227

4846.3 9 Violet Pseudo-Continuum for H$

4861.3 18 H8

4876.3 9 Red Pseudo-Continuum for H8

4948.0 15 Violet Pseudo-Continuum for TiO 4954

4964.0 15 TiO 4954 (0,0) Bandhead Table 8. Spectral Indices Measured in the Spectra of Non-Degenerate CPM Stars.

Name n S(CaI)a S(Ti0)a C(HK) a S(HK) a S(B) a AS(HK) AS (3)

G272-B5A 2 1009 1 970 10 440 1 227 3 867 2 11 -11 -31°1454 2 925 6 974 10 343 17 240 3 921 0 59 -6 LP888-25 2 603: 0 972: 39 885:126 197: 22 1041: 1 -62: 58: 02 Eri A 2 940 2 1016 6 710 1 184 3 926 1 1 7 02 Eri C 2 512 2 598 5 766 26 511 3 1113 2 189 120 LP891-13 2 670 3 719 5 910 24 819 37 1512 36 590 537 +12°937 3 984 3 987 6 464 9 232 3 885 7 19 4 -59°1275 2 985 6 992 12 493 6 198 0 916 5 -9 30 -28°3361 3 886 7 983 16 889 31 196 1 947 7 16 12 -18°2482 2 846 6 971 2 1067 2 179 2 979 11 -3 37 G114-B8A 3 952 6 973 3 408 2 262 2 917 1 55 31 G114-B8B 2 834 9 973 19 821 64 150 5 948 3 -30 13 -31°7352 2 829 3 963 1 909 8 153 2 973 2 -29 33 LP902-30 2 538 8 674 6 884 54 244 8 976 14 -12 -6 L753-40 2 599 1 809 5 808 3 227 1 1025 21 -35 41 L753-39 2 533 9 781 2 829 34 224 9 995 10 -76 5 -18°3019 1 624 854 793 248 1026 7 48 G163-B9A 2 • 962 1 956 0 387 4 297 2 936 1 96 44 L970-27 1 556: 693: 834: 296: 962: 22: -24: -25°8487 2 931 4 990 4 519 7 217 1 928 3 13 39 L1046-18AB 1 567: 694: 843: 339: 989: 32: -2: +23°2539 3 978 5 984 6 742 32 229 5 939 2 46 20 W485B 1 517: 625: 899: 342: 932: 20: -61: LP498-25 1 586: 754: 1034: 281: 967: 13: -18: L619-51 1 603 831 864 292 1006 33 23 LP740-11 1 927 978 422 225 922 5 47 LP740-12 1 926 992 308 266 883 23 25 LP916-26 1 546: 622: 889: 365: 955: 72: -34: LP684-1 2 747 1 928 2 904 37 213 3 1016 10 12 54 LP684-2 1 591: 747: 918: 254: 1014: -14: 29: -38°10983 2 951 1 979 3 508 29 243 0 943 1 43 49 L266-195 2 701 0 926 10 838 8 302 8 1017 4 85 47 -8°4352AB 2 562 1 809 37 800 8 416 9 1054 1 116 64 -8°4352C 2 503: 5 690: 24 978:83 286: 2 1023: 39 -51: 28: G154-B5A 2 669 15 919 13 853 26 307 2 1002 5 56 21 +9°3501 1 1019 954 910 211 957 30 21 -4°4636A 1 944 972 547 208 955 17 51 -4°4636B 1 957 1014 647 207 976 26 40 LDS678B 1 496: 721: 881: 320: 1013: 9: 21: L1142-88 1 704 701 802 316 966 36 -21 R193 2 515 1 709 8 767 3 369 2 1048 6 104 64 -0°4234 2 884 5 950 0 684 22 291 1 1008 3 110 70 -51°13128 1 552 757 604 645 1200 374 215 LDS766B 2 556::10 617: 12 893:114 371: 1 957: 22 91: -30: LDS785B 1 637: 832: 813: 299: 985: 56: 6: Table 8 (Continued).

Name n S(CaI)a S(Ti0)a C(HK) a S(HK) a S(6) a AS (HK) AS(8)

LP581-36 2 640 1 847 7 863 33 300 0 1000 1 46 18 G275-B8A 1 979: 1080: 445: 310: 892: 94: 14: G275-B8B 2 889:10 864: 1 924:77 261:18 985: 5 80: 47: LP933-66 1 515: 675: 709: 281: 1028: -33: 36: LP935-14 2 697 15 882 1 859 36 258 4 1012 19 39 41 G273-B1A 2 982 1 955 0 740 4 231 2 953 8 31 59 LDS826B 1 560 682 957 313 960 30 -28 G275-B17A 1 999 967 455 279 920 65 40 Table 9. Spectral Indices Measured in the Spectra of Bright F-M Dwarfs.

HD n Sp B-V S(Cal) a S(TiO) a C(HK) a S(HK) a S(8) a AS(HK) AS(B) f (h k :

4628 1 K4V +0.88 911 991 864 208 933 28 5 200 9562 2 G2V +0.65 1021 11 976 3 538 17 250 8 912 1 53 15 148 16673 3 F7V +0.52 991: 1 981: 3 409: 10 226: 7 862: 6 8: -14: 203 17925 2 KOV +0.87 970: 6 991: 10 876: 10 197: 2 941: 1 16: 14: 500: 22049 1 K2V +0.88 907 961 785 194 928 14 0 350: 23249 1 KOIV +0.92 1039 944 772 175 921 -5 -12 156 26965 2 K1V +0.81 940 2 1016 6 710 2 184 3 926 1 1 7 230 30495 2 G1V +0.60 963 3 976 15 539 3 195 1 883 9 -9 -6 230 32147 2 K5V +1.06 925: 14 971: 5 1029: 29 186: 4 946: 8 0: -3: 300 45067 2 F8V +0.56 995 4 973 5 456 3 244 2 910 4 33 27 141 76151 3 G3V +0.65 1005: 5 986: 7 577: 64 222: 18 916: 8 25: 19: 260: 81809 2 G2V +0.64 999: 1 979: 2 517: 9 187: 3 922: 8 -11: 27: 160 88737 3 F5V +0.63: 1025 6 980 3 544 7 246 1 894 1 46 0 215 AD Leo 1 dM4-5e +1.54: 542 847 977 602 1140 322 153 - 95735 2 M2V +1.51 644: 8 878: 0 1173: 35 281: 9 993: 3 10: 8: 520: 106516 2 F6V +0.47 978 2 979 8 338 5 281 3 918 6 52 50 207 109379 1 G5III +0.89 1138 984 578 203 978 23 49 - 115383 2 F8V +0.59 1022: 7 1011: 6 546: 14 211: 2 888: 13 5: 0: 250: 115404 3 K3V +0.95 913 8 972 16 809 11 249 1 958 4 68 22 480: 126053 1 G3V +0.64 925: 1039: 524: 219: 935: 21: 40: 165: 136202 1 F8V +0.54 990 938 460 249 883 35 3 139 155885 3 K1V +0.86: 932 13 966 21 818 51 252 6 975 11 71 49 360 155886 3 K1V +0.86: 926 9 965 9 847 54 250 9 967 2 69 41 425 156026 1 K5V +1.16 801: 977: 1100: 247: 961: 51: 2: 500: 190406 2 G1V +0.60 996 0 1013 5 588 19 266 6 935 5 62 46 187 206860 1 GOV +0.59: 1005 993 509 267 927 61 39 250 88

This dependence was readily evident from a visual examination of the

vidicon spectra and suggested that measurements of the strength of this

feature could provide a fairly accurate means of objectively classify­

ing dK and dM stars among the CPM RD's. The red wing of this feature

is severely contaminated by lines of Crl and the nearby diffuse G-band

of CH so that accurate measures of the equivalent width of Cal 4227

cannot be obtained. The violet wing is relatively uncontaminated.

Thus, following the example set by Equation III-4, a normalized flux

ratio which measures the strength of this feature can be defined:

S(Cal) = FCaI x 1000 V-l

where is the integrated flux in a 30 X window centered on Cal 4227 and Fy is the integrated pseudo-continuum flux in the 30 8 window given in Table 7. With the normalization constant given in Equation

V-l S(Cal) would be expected to vary between 0 for a saturated line and 1000 for a pure continuum region.

Figure 21 illustrates the strong dependence of S(Cal) upon spectral type for stars in Tables 8 and 9 which have previously published spectral classifications. S(Cal) appears to vary linearly with spectral type for faint dK and dM stars among the vidicon spectra.

The bright stars from Table 9 (large symbols) show a much tighter correlation perhaps because most have spectral classifications supplied by Wilson (1978) on the MK system (Morgan et al., 1943). Spectral types for the fainter CPM RD's from Table 8 (small symbols) come from various sources which in several stars disagree by as much as two or three spectral subclasses. Three objects lie significantly off HD23249 1000

900 S(Cal)

800

G24-10 700

600

• HD at • CPM ID 500 a> K0 K1 K2 K3 K4 K5 K7 MO Ml M2 M3 M4 M5 SPEC. TYPE

Figure 21. The Dependence of S(Cal) on Spectral Type in K-M Dwarfs.

00 VO 90

the mean relation. HD23249 has been classified by Keenan and Pitts

(1980) as a subgiant and falls noticeably above the general correlation

indicated. BD+9°3501 has anomalous UBV colors and may be composite.

G24-10 was noted as having a very close optical companion during the

vidicon observations.

Since spectral classes form essentially a temperature sequence

it was expected that S(Cal) could provide a means of estimating colors

for the CPM RD's which lack photoelectric photometry. Figure 22

displays the observed correlation between S(Cal) and B-V color for all

objects in Tables 8 and 9 which have prior UBV data. A regression

line obtained from all late type dwarfs in this figure yields the

empirical relation

S(Cal) = 1390 - 536 (B-V) V-2

for B-V >+0.80 (region indicated by arrow in Figure 22). Fits using

only the bright stars from Table 9 do not appreciably alter this

result. From the residuals of Equation V-2 it appears that S(Cal) is

capable of estimating B-V colors for faint dK and dM stars to an accur­

acy of aB_y~ iO.O?111. The G5 III star HD109379 appears well above the value of S(Cal) implied for zero absorption in Figure 22 . it thus appears that the cause of the weak luminosity dependence mentioned in the discussion of Figure 21 lies in the contamination of the pseudo­ continuum band which defines F^ in Equation V-l among higher luminosity stars. Fortunately nearly all the nondegenerate CPM stars appear to be dwarfs.

Fifteen CPM RD's have R,I photometry reported by Eggen (1979a).

From these observations S(Cal) was found to exhibit a strong though 91

H 0 109379 1100

HD 33349 1000 S (Ca I) 900

800

700 G 34 -10

• \ 600

• HO M 500 • CPM «D

+0.6 +0.8 +1.0 +1.2 +1.4 +1.6

Figure 22. The Dependence of S(Cal) on B-V Color. 92 non-linear empirical dependence on R-I color

S(Cal) = 1280 - 1180 (R-I) + 453 (R-I)2 V-3 for R-I ^+0.4, which permits estimation of R-I colors in K and M dwarfs to an accuracy of - ±0 .10m from the vidicon spectra.

Keenan and McNeil (1976) note the enhancement of Cal 4227 in subdwarf spectra. Although the Hg vs. Mg diagram (see Figure 3) indicates that few if any CPM RD's are kinematically true subdwarfs a means of testing this assumption exists if the spectral classes and/or photometric colors derived from Cal 4227 can b§ compared to independent measurements of these quantities. In subdwarf spectra

S(Cal) should predict a later spectral class and redder color than actually exists. Several objects in Figure 21 have published spec­ tral classes earlier than those estimated from S(Cal). However these objects lie much closer to the S(Cal) vs. B-V relation indicated in

Figure 22 . We conclude that the discrepancies existing in Figure 21 are most likely due to errors inherent in the subjective classification of photographic spectra and not to the subdwarf nature of these objects.

It appears that none of the dK and dM stars having published UBV data are true subdwarfs. The possibility exists however that some subdwarfs exist among the objects lacking photometric data. Thus an independent check on the spectral classes and colors derived from Cal is desirable.

The appearance of the TiO 4954 (0,0) bandhead helps define spec­ tral subclasses among M stars (cf. Morgan et al., 1943, and Keenan and

McNeil, 1976). Liebert et al. (1978) point out the potential usefulness of this feature for optical spectral classification of late dM stars 93

using MC detectors. We define a normalized flux ratio similar to S(Cal)

which measures the strength of the depression caused by the TiO 4954

bandhead

FTi0 S(TiO) = x 1000. V-4 FV FTiO anc* FV are ^nte8rate^ fluxes in 30 X windows located on either side of the bandhead respectively (see Table 7 ). S(TiO) was found

to correlate well with R-I color. The empirical relation

S(TiO) = 1020 - 49.9 (R-I) - 193 (R-I)2, V-5

valid for R-I ^+0.4, allowed estimates of R-I to be determined for the

dK and dM stars which lack R-I photometry to an accuracy of ~

±0.09m (based on the residuals of objects having R,I data). These R-I

estimates were compared to those predicted by S(Cal) in Equation V-3.

Among the objects lacking independent photometric data no discrepancies

which could be ascribed to an enhancement of Cal 4227 relative to TiO

4954 were found. We thus conclude that it is unlikely any true sub­

dwarfs exist among this subset of CPMLIST RD’s.

Quantitative Classification of dF and dG Stars: The Call H&K Continuum Index

Although Cal 4227 and/or TIO 4954 allow reasonably accurate estimates of spectral type and photometric color to be obtained for late type dwarfs neither of these features is strong enough to be of use in stars earlier than late G or early K spectral class. Many of

the CPM RD’s which lack observational material prior to the vidicon 94

spectra appeared to be F-G dwarfs. Consequently a means of objectively

classifying and estimating colors for these objects was desirable.

Vaughan and Preston (1980) have shown that a color index formed

from integrated pseudo-continuum fluxes situated on either side of

Call H&K correlates well with photometric color for stars of B-V <^+1.1.

Following Equation I1I-5 we define such a color index for the vidicon

Fv C(HK) = -2500 log — V-6 FR

where F^ and FR are pseudo-continuum fluxes integrated over the 20 8.

windows given in Table 7 . As is the case for S(Cal) and S(TiO) the

normalization constant is used to provide a convenient format in

Tables 8 and 9.

Figure 23 displays the roughly linear correlation found to exist

between C(HK) and spectral type for stars with published spectral clas­

sifications. The bright stars in Table 9 (large symbols) again fall

closest to the mean relation.

The correlation between C(HK) and B-V is plotted in Figure 24.

A formal least squares fit yields the relation

C(HK) = -134 + 1080 (B-V) V-7

for objects with B-V <_ +1.1 (arrow in Figure 24 ). From the residuals it appears C(HK) predicts B-V to an accuracy of ±0.04m . HD23249 a B—V * (K0 IV) and HD109379 (G5 III) have been excluded from the calibration since C(HK) appears to be luminosity sensitive. Evidently C(HK) is of little use for stars later than about K5 (B-V > +1.1). An overlap of nearly one full spectral class (+0.8 B-V ± +1.1) exists between the useful limits of C(HK) and S(Cal). Spectral types and B-V estimates 1200

1000 AD Lao C(HK)

800

600 HD109379

• HD 400 • CPM RD

F6 F8 GO G2 G4 G6 G8 KO K2 K4 MO M3 M7 SPEC. TYPE

Figure 23. The Dependence of C(HK) on Spectral Type in F-K Dwarfs. 1200

1000

800 HD23249

600 • HO 109379

400 • HD • C P M RD

+0.6 +0.8 +1.0 +1.2 +1.4 +1.6 B-V

Figure 24. The Dependence of C(HK) on B-V Color. 97 listed in Table 8 for objects in this range are averages predicted by

C(HK) and S(Cal), which generally agreed to within one or two subclasses

( A(b-V) = ±0.07™).

It appears from this study that with some caution MC measurements of Cal 4227 and Call H&K pseudo-continuum fluxes may be used to deter­ mine medium quality estimates of spectral type and photometric B-V color for faint late type dwarfs from blue region vidicon spectra alone. In the absence of red or infrared spectra or photoelectric photometry such estimates are probably at least as accurate, require less observation and data reduction time, and yield more objective results than more conventional photographic methods.

Chromospheric Activity in CPM RD's: Call H&K and HB Emission

The strengths of the chromospheric emission lines of Call H2 and

K2 (henceforth abbreviated HKe) at the centers of the photospheric absorption lines in bright F-M dwarfs have been shown to be strong indicators of stellar age (Wilson, 1963) and useful in the detection of analogs to the solar activity cycle and transient chromospheric events in such stars (Wilson, 1978 and Baliunas et al., 1981, respectively).

Several investigators (cf. Wilson, 1968, Vaughan et al., 1978, and

Barry et al., 1981) have constructed moderately high resolution instru­ mental systems capable of measuring HKe only in bright stars, since such emission generally exhibits FWHM K 1 X. The results of a pilot

* program undertaken in 1978 at Lowell Observatory with the Perkins 1.8m telescope and Boler and Chivens photographic image tube spectrograph indicated that HKe levels strong enough to measure even at relatively low resolutions ( '''lO &) exist among some faint CPM RD's. To determine 98 the nature of such emission and possibly extend some of Wilson's results to much fainter stars than heretofore attempted the vidicon spectra have been evaluated for the presence of HKe.

From Equation III-4 we define a flux ratio for HKe

S(HK) = - 7 -;- V-8 V R where, as indicated in Table 7, F ^ and F ^ are the integrated fluxes in 9 8 windows centered on Call K2 and H2 respectively, while F^ and F^ are the integrated pseudo-continuum fluxes previously defined for the color index C(HK) in Equation V-6 . The sizes ‘of the F ^ and F ^ win­ dows are fixed by the instrumental resolution and therefore admit considerable amounts of photospheric flux which contaminate the HKe measurements. It should be also noted that F and F„ are not true V K continuum points due to extensive line blending which exists in the blue-violet spectra of late type" stars. Both of these effects are functions of spectral type and must be considered if emission levels for stars of different temperatures are to be intercompared. A summary of the values of S(HK) measured in the vidicon spectra appears in

Tables 8 and 9.

Haro (1976) points out that strong HKe and Balmer emission appear to be associated with flare activity among young members of the Pleiades and Orion Nebula region. Flare stars have been shown to be less common among members of older clusters though measurable amounts of HKe and

Balmer emission often persist in objects which show infrequent outbursts or none at all, implying that chromospheric activity decays very slowly.

Further supporting this are the results of Skumanich (1972), Blanco et al. (1974), and Barry et al. (1981) which indicate that HKe persists 99 in stars at least as old as the sun. Balmer emission might also be expected to be a useful index of stellar age. Haro (1976) does however mention that more young stars appear to exhibit HKe than Balmer emis­ sion, implying the timescale for the decay of such emission is shorter than for HKe.

Although the vidicon spectra do not extend to Ha we define a normalized flux ratio to measure Hg emission (Hge) among the CPM RD's

Fg s(e) H v + v x 1000 v_9 rv R where FD is the integrated flux within an 18 R window centered on Hg p and Fy and FR are integrated pseudo-continuum fluxes on either side of Hg . In principle S(g) = 1000 for zero absorption/emission. The

S(g) observations obtained from the vidicon spectra of late type dwarfs appear in Tables 8 and 9.

It is of interest to determine the level at which HKe and Hge become detectable in the vidicon spectra. Nearly all the bright F-M dwarfs in Table 9 have been selected from Wilson's (1978) compilation of long-term high resolution observations of HKe. Wilson has identi­ fied many objects during the course of this program which undergo stochastic and/or cyclic changes in HKe implying that analogs to solar activity exist in other late type stars. Stars from this list which were accessible from the southern hemisphere where the vidicon spectra were obtained and which Wilson has shown exhibit constant or quasi-periodic changes in HKe since 1968 when the program began allowed rough estimates for the HKe intensities for the 1979-80 observing seasons to be made. Unfortunately these objects span less than half 100

the range in HKe intensities exhibited by the northern hemisphere stars on Wilson's program.

Figure 25a illustrates the relation between S(HK) and the esti­ mated 1980 values of F(HK), Wilson’s instrumental HKe measurements.

Stars of spectral types F, G, and K-M are indicated by different symbols.

From this figure it is apparent that only for K-M stars does there appear to be a correlation between S(HK) and F(HK)^ggQ. A formal least squares fit for these stars (solid line) indicates that S(HK) is about 20% as sensitive as F(HK) to the presence of HKe among K-M dwarfs, somewhat better than the ratio of the resolutions of the two instrumental systems ( ^ X for the vidicon, X for Wilson's HKP-1 spectrophoto­ meter) . The fit does not change appreciably when Wilson’s latest 1977 observations are used (dashed line). Error bars for F(HK) are taken from the standard errors determined by Wilson from fluctuations observed in these stars since 1968. One of the least active dM stars which shows strong HKe, HD95735 (M3 V), is known to occasionally exhibit strong outbursts (Wilson, 1978). The maximum and minimum HKe levels

Wilson has reported for this star are indicated (arrows).

Figure 25b documents a correlation which appears to exist between

H0e and HKe among the vidicon spectra of Wilson's bright F-M dwarfs.

Surprisingly the F and G stars appear closer to the relation suggested by the K and M stars than was the case in Figure 25a. This may be a result of the fact that H 6 occurs in a much less complex region of the spectrum of late type stars than Call H&K. At any rate the cor­ relation between S (8 ) and F(HK) does not change appreciably when

Wilson's last reported 1977 values of F(HK) are used (dashed line) in 101

KD95735 280

KH 240 S(HK)

200 H a-t

160

1000 H D 95735 O G • K-M

960

S(P)

920

880

200 400 500300 600

F(HK) 1980 «st. Figure 25. Comparison of (a) S(HK) and (b) S(B) Vidicon Indices With Call H&K Fluxes Reported by Wilson (1978). 102

the fits instead of the 1980 estimates (solid line).

Figure 25b suggests that S(3) should correlate well with S(HK)

since these features were observed simultaneously with the same instru­ mental system. Presented in Figure 26 is further evidence of such a

correlation among the bright F-M dwarfs in Table 9 . In this figure

K and M stars form a tight, apparently linear sequence

S( 6 ) = 0.641 S(HK) + 807. V-10

The F and G stars, though suggesting a generally parallel positive correlation between S( 3) and S(HK) are clearly well below the relation of Equation V-10. Note the anomalous position of HD109379 (G5 III).

He contaminates the absorption profile of Call H. In order to test whether the apparent correlation between S( 8 ) and S(HK) in Figure

26 is due to the effects of He on the S(HK) index a plot of S( 8 ) vs. S(K) was constructed, S(K) being a flux ratio for Call K. The relation which resulted was nearly indentical to Equation V-10, indi­ cating that no significant contamination of S(HK) by He occurs among these spectra.

Figures 25 and 26 imply that S(HK) and S( 8 ) are sensitive to a star's color as. well as emission level. The most likely cause is the contamination of the Call H&K feature bandpasses by photospheric flux, a problem previously noted. Because a deeper line relative to the pseudo-continuum can result from a hotter star as well as an older star the effect of temperature should be calibrated out. The procedure used here is illustrated in Figures 27 and 28 where the values of

S(HK) and S( 8 ) from Tables 8 and 9 are plotted against published

B-V colors for these objects or estimated from S(Cal) and/or C(HK) by 103

1000

SO)

950

900

• G • K-M 850 150 200 250 S (HK)

Figure 26. The Correlation Between S(B) and S(HK) Among Bright F-M Dwarfs Observed With the SIT Vidicon. 800 • HD • CPM RD • CPM RD,B-Vestd 700

600 S(HK) 500

400

300 V47I Tau <3* OUT • • i 200 O, Erl A 0©

+0.6 +0.8 +1.0 +1.2 +1.4 +1.6 B-v Figure 27. The Dependence of S(HK) on B-V Color. 105

1500 • HD • CPM RD

1400 • CPM RD,B-V Esro

1300 S(« 1200

• ad U o 1100 ©,ErlC®

1000

900 O, Erl A

+0.6 +0.8 +1.0 +1.2 +1.4 +16 B-V Figure 28. The Dependence of S((3) on B-V Color. 106 the methods outlined earlier in this chapter. The general appearance

of these plots is reminiscent of the relation between F(HK) and b-y

found by Wilson (1968) except that the vidicon observations extend to much redder and fainter stars.

The lower envelope of points in each diagram is defined by fitting an empirical curve through the two or three lowest emission stars within each 0.1m interval in B-V which have published UBV data. These curves provide a relative "zero emmission level" for HKe and HBe against which other stars may be compared. It should be noted that such a procedure requires a large sample of stars to define the lower envelope and that the stars for which B-V colors were estimated have not been included in the fitting procedure, which in any case provides only a rough way of calibrating out the color dependence of the S(HK) and S( 6 ) measure­ ments. Barry et al. (1981) have used a similar technique to calibrate

HKe levels among F and G dwarf cluster members.

Relative emission level indices AS(HK) and As( 6) can now be defined such that

4S(J) H S(l)observed - S(J)0 V-ll where S( ^)0bservecj represents the S(HK) and S(8 ) indices measured from the vidicon spectra and S(1 )q represents the "zero emission level" at a particular B-V indicated by the curves in Figures 27 and 28 . The computed values of AS(HK) and As(B ) are recorded in Tables 8 and 9.

The correlation between AS(B ) and AS(HK) for all non-degenerate stars observed with the vidicon is indicated in Figure 29 by a solid line. F, G, and K-M stars have again been keyed by different symbols.

The bright HKe standards previously plotted in Figure 26 lie within the dashed box at the lower left of Figure 29 . We note several LB91-13 500 ~ • K-M Low Wt. 400 A S W 300 -

//•-irnm 200 -

AD L«o

-0 4 2 3 4 100 - ■193^; • 4352 A / i out V4T1 Tau

0 100 200 300 400 500 600 AS(HK)

Figure 29. The Correlation Between AS(3) and AS(HK) Among F-M CPM Dwarfs. interesting features of this diagram. First, there no longer appears

to be a separation between stars of F-G and K-M spectral types which existed in Figure 26 , hence it appears that color effects have been successfully removed. Second, it is apparent that nearly all stars in the present sample exhibit relatively low emission levels as would be expected for CPM systems containing WD components which may have Q ages > 10 y. Third, the low emission end of the AS(6) vs. AS(HK) sequence (lower left in Figure 29 ) appears to be bifurcated. Whether this is a result of some difficulty in the calibration scheme or sup­ portive evidence of an apparent deficiency of stars with intermediate

HKe levels in the solar neighborhood (Vaughan and Preston, 1980) cannot be determined from the present low-resolution sample. We do note, however, that internal standard errors in AS(8) and AS(HK) are gener­ ally smaller than the symbol sizes in this figure. Error bars for objects with larger uncertainties have been plotted. Finally we note that the four objects with highest emission levels in this sample are all well-known flare stars and/or unresolved binaries— both character­ istics often associated with HKe and H8e. As further evidence that HKe and HBe are being detected loci of the Hyades WD eclipsing binary V471

Tau (Nelson and Young, 1970), for which vidicon observations inside and outside eclipse of the WD component were obtained, are plotted in Figure

29 . This system has been shown to exhibit moderately strong HKe outside eclipse and variable H8 flux (Oswalt, 1979).

Wilson (1963) first pointed out that comparable HKe strengths exist among members of bright visual binaries. Since HKe and, because it appears to be correlated with HKe, H6e decline with stellar age the 109 coevality of the 11 RD + RD CPM pairs in Tables 8 and 9 may be tested by comparing the levels of emission exhibited by the components.

Figures 30 and 31 compare the relative emission levels AS(HK) and

A S(B) between the components of these systems. Seven of eleven pairs in Figure 30 exhibit comparable AS(HK). The exceptions include:

O2 Eri A/C: Component C is a well-known flare star with exceptionally strong emission.

-31°1454 A/B: Spectrum of the faint companion is very weak. AS(HK) and AS( 8 ) should not be considered reliable.

G114-B8: A Giclas pair of uncertain CPM. The magnitude difference is too great for a pair of MS stars.

-8°4352 AB/C: The primary consists of a close pair 0.2” apart which could not be resolved with the vidicon.

In Figure 31 the same four stars are discordant. Most of the remaining primaries appear to have stronger levels of H8 emission than their companions (i.e. fall to the right of a line of unit slope in the figure). Whether this effect is real or an artifact of the calibra­ tion process is unclear and will require higher dispersion spectra to clarify. The components of HD155886/5, a bright visual binary observed by Wilson (1978), appear to exhibit comparable emission levels in both diagrams. In general it appears that most of the components of CPM RD pairs exhibit HKe levels of comparable strength. Although the results for HBe are less certain a weak positive correlation is suggested.

Skumanich (1972) and Blanco et al.(1974) have shown that the strengths of Call H&K emission decay with stellar age according to an inverse square root law which appears valid to at least an age comparable to that of the sun. Since the temperature of a WD depends upon the time it has been losing thermal energy it might be expected that an iue3. ComparisonofAS(HK) Exhibited Componentsby ofCPM Figure30. CO X <

SECONDARY 50 -5 100 150 50 Binaries ContainingBinaries TwoF-M Dwarfs. / #,*,• / i / T G G FG + KM G G O FG ♦ FC LowWt. KM+KM f 5 100 50 0 2EriA/C o • o / * AS(HK),

D5865 ✓ YG275-B8 ✓ T HD155886/5 V LP740-12/11 / / / i G114-B8 -31*1454 A/B PRIMARY A -8*4352 AB/C / / / / / / no Figure o t <

SECONDARY 1 Comparisonof AS(3) ExhibitedComponentsby of CPM Binaries 31. 100 25 50 75 0 Containing TwoF-M Dwarfs. -31*1454A/B O O O FG + KM • • KM+KM FC+FG 2 50 25 0 •o, •o,

LP740-12/11 AS(0) Eri G114-B8 / A/C HD1558B6/5/ PRIMARY G275-B8 H— M -8°4352 AB/C / r - 75 111 empirical relation exists between the temperatures of a WD and the HKe

strength of a RD companion. Figure 32a exhibits the observed relation

between AS(HK) for the nondegenerate stars and log for the WD compan­

ions having well-determined temperatures by Koester et al. (1979) and/or

Shipman (1979). Some systems which display exceptionally high levels

of emission are apparent. The anomalous emission level of o3 Eri C

(dM6e) has been noted earlier. Joy and Abt (1974) note that essentially

all dwarfs later than dM5.5 are dMe stars which exhibit strong HKe and

HBe. Such objects are likely to have completely convective envelopes

(cf. Mullan, 1975) which may be expected to maintain chromospheric

activity well beyond the timescales for decay found among earlier spec­

tral types. Thus it is not surprising that o^ Eri A (dKl) exhibits HKe

levels more comparable to the main grouping of CPM pairs in Figure 32a.

G154-B5 is a Giclas pair of uncertain CPM. The system -51°13128/L283-7

(AS(HK) = 215, offscale, indicated by an arrow) has a primary with very

strong, doubled emission lines of Call H&K and Hg and hence appears to

be an unresolved binary. It is interesting that this system has a

degenerate companion which has some of the strongest carbon bands

exhibited by C2 WD's (Wegner, 1973). Eggen and Greenstein (1967)

resolve the apparent "paradox" of a strong emission star like R193

paired with the cool and presumably old DC WD VBsll by concluding that

the latter is a very low mass degenerate star. Such objects cool

rapidly enough to explain why emission levels in R193 (dM6e) have not not completely decayed. As for the remaining CPM pairs it is apparent

that little if any correlation exists between A S ^ K ) ^ and log T3 ^

even when the peculiar systems just discussed are discounted. We 113

200 i ) ' -•1*13128 .f A/B f O. Eri B/C

150 a s (h k )rd

x © 100 £ R193/VB.11

G154-B5 I 50

O 0 £ O j Ert A/B

Er! B/C 100 (b) -51*13128 A/B

IB193/VB*11 50

AS (fi)n G154-B5 I J J Oj Eri A/B 0

• DA o • DB * O DC -50 •,0 , 0 Low Wt.

I______I______I______I______I_ 3.5 4.0 4.5

log T3 wd

Figure 32. Comparison of (a) AS(HK) and (b) AS(6) of F-M CPM Dwarfs to Log of WD Companions. 114 surmise that the effects of the intrinsic mass dispersion and/or MS lifetimes of the WD progenitors have a large influence on the perceived relation (or lack thereof!).

Figure 32b displays the relation between AS(S)jyj and log yjj.

Although CPM pairs containing DA WD's appear distinguishable from those which contain non-DA degenerate stars, as in Figure 32a no clear-cut correlation is evident.

Summary and Conclusions

The purpose of this investigation of the spectroscopic features common to late type members of CPM binaries was to see if relatively low resolution spectra could provide clues to the relationships which have been presumed to exist between the companions. In particular the following points arose:

1.) Quantitative spectral classification of F, G, K, and M dwarfs have been shown to be obtainable from the vidicon spectra which avoid many of the difficulties in subjectively ascribing spectral types to faint stars from photographic spectra. Furthermore blue region spectra alone appear to be sufficient to establish fairly accurate estimates of photometric color for such objects which, for many purposes, may avoid the necessity of large investments of observing time in sup­ plementary photometry. Among the CPM RD's none appear to be subdwarfs.

2.) Quantitative measurements of Call H&K emission, known to correlate with stellar age, can be obtained at significantly lower resolutions than heretofore attempted. In principle this may permit future work on the HKe vs. age calibration begun by Skumanich (1972) to be extended to very faint cluster members and general field stars. 115

3.) Apparently for the first time a quantitative relation between

the relative levels of Call H&K emission (HKe) and Hg emission (Hge)

has been demonstrated. Although Hge does not appear to be as sensitive

to chromospheric activity as HKe, it is located much nearer to the

energy maxima of late type stars and is less severely contaminated by

nearby spectral features. In very faint stars these advantages may be

important. We infer that measurements of Ha emission may prove to be

significantly better correlated to HKe.

4.) The relative emission levels of the components of CPM red

dwarf pairs appear to be roughly correlated. Combined with the kine- matical evidence these results support the supposition that most are

true physical pairs and share a common origin. Discordant emission

levels among such pairs seem to be relatively good indicators of unre­ solved binary components, stars subject to stochastic flare activity, and, in several cases where weak kinematical evidence exists for true physical association, cast further doubt on the CPM nature of the component stars.

5.) Nearly all RD stars with a WD companion exhibit very low levels of HKe and Hge, implying ages qualitatively commensurate with cooling timescales and progenitor lifetimes of the degenerate components.

Although a quantitative relationship between the emission levels of RD components of CPM pairs and the temperatures"of the WD companions does not appear evident among the low resolution vidicon data the subject remains fertile ground for future work involving an absolute calibration of HKe and/or Hge for such systems. CHAPTER VI

SUMMARY

This study was designed to investigate the spectroscopic properties

of faint (m > 14) CPM binaries which contain WD components. Such work P S would not have been feasible prior to the development of high quantum

efficiency MC detectors like the CTIO SIT vidicon. Until now such

equipment has not been extensively employed in the analysis of absorp­

tion line features in faint stars. Consequently a major result of this

investigation has been to demonstrate that MC spectra are well-suited

to the analysis of degenerate and non-degenerate stellar spectra because

of the increased sensitivity, linear response, better spectrophotometric

accuracy, and greatly simplified reductions afforded by digital spectra.

In addition, with an adequate determination of nightly extinction and

flux calibrations, integrated spectra may also provide magnitudes and

colors of accuracy approaching that of conventional filter photometry

while retaining full spectral information.

This observational survey was based on CPM binaries selected from

the Luyten (1963, et seq.) and Giclas et al. (1971, 1978) proper motion

surveys. The crude magnitude and color estimates supplied by these

catalogs have been shown to allow adequate criteria to be defined for

the selection of degenerate suspects among the CPM binaries. This

investigation has resulted in new spectroscopic identifications of roughly one dozen WD's.

116 117

Among the CPM binaries containing a WD and a late type MS companion

Luyten (1974) has pointed out specific relationships between the magni­

tudes and color classes exhibited by the component stars. This study

has confirmed the existence of at least two groups of such binaries

using UBV data which are now available for many of these stars. These

groups are apparently a result of the intrinsic shape of the WD cooling

sequence and the downturn in the cool end of the MS seen in the H-R

diagram as well as observational biases inherent in previous searches

for WD's among the CPM binaries.

We have shown that reasonably accurate temperatures and gravities

for DA and DB WD’s can be determined from straightforward measurements of absorption line strengths and widths. With grids of model atmo­ spheres which are now available (cf. Koester et al., 1979; and Koester,

1980) this method may in many cases avoid the necessity of computing atmospheric models for individual WD's.

The present study has greatly expanded the spectroscopic information available for the non-degenerate members of CPM binaries. The vidicon spectra have been used to objectively classify these stars using simple measurements of Call H&K, Cal 4227, and TiO 4954. These measurements have also been shown to provide a means of estimating fairly accurate

B-V and/or R-I colors, avoiding supplemental photometry. An important result of this portion of the study is that few if any of the CPM stars in the present sample appear to be, kinematically or spectroscopically, true subdwarfs. Most are high velocity disk objects. This result provides an important clue to the WD birthrate in the Galaxy. 118

We have sought to test the assumed coevality and physical associ­

ation of the CPM binaries which do not contain WD's by using Wilson's

(1963) discovery that Call H&K emission strengths of visual binary

components are comparable. Since such emission is correlated to stellar age (cf. Skumanich, 1972) the components of wide non-degenerate CPM pairs should exhibit comparable emission levels. In most cases this proved true although discordant emission levels were observed in a few pairs whose physical association was uncertain for kinematical reasons.

Hence the technique provides a secondary check where CPM is uncertain.

Apparently for the first time it has been demonstrated that Hg emission correlates well with Call H&K emission among late type stars.

This result is somewhat surprising in light of Haro's (1976) observation that Hg is not always in emission among (young) flare stars which exhibit strong Call H&K emission. Possibly the correlation improves as emission levels decline with stellar age since nearly all CPM stars in the present sample display relatively low levels of Call H&K and Hg emission. Hg is located closer to the energy maxima and is therefore easier to measure in faint late type dwarfs. The results for Hg suggest that Ha emission measurements may also prove useful in the calibration of stellar ages and the study of chromospheric activity.

Although the resolution of the CTIO SIT vidicon spectra is insuf­ ficient to attempt an absolute Call H&K (or Hg) emission vs. age cali­ bration, recent improvements in this system (cf. Hesser and Harris, 1981) admit the possibility that such work may soon be undertaken. Among faint open cluster dwarfs Barry et al. (1981) have attempted this with similar equipment. 119

We anticipate that radial velocities (cf. Hesser and Harris, 1981)

will soon be obtainable from vidicon spectra. This would provide a

means of deriving the intrinsic gravitational redshifts of faint CPM

WD's since the radial velocity of the non-degenerate companion is

essentially the systemic velocity (orbital velocities are in general

< 1 km/sec in CPM pairs). Hence the difference between the radial

velocities of the two components gives the gravitational redshift of

the WD directly.

The final portion of this investigation compared the (age-dependent)

Call H&K and Hf5 emission strengths of non-degenerate CPM stars with

the (cooling time-dependent) temperatures of WD companions. Although

no clear-cut'Correlation was found these results remain inconclusive.

Future work will require higher resolution and corrections for presumed

WD masses and pre-WD lifetimes.

The primary objective of this investigation has been to call atten­

tion to the fact that CPM pairs provide an important (albeit unglamorous) means of obtaining fundamental properties of both degenerate and non­ degenerate components. If any one statement encapsulates the spirit of this study Luyten (1969) has come closest:

"Observationally speaking, therefore, these proper-motion doubles are the most common type of binaries in space. Nevertheless, double star observers and theoreticians alike continue to ignore them, to sweep them under the rug, so to speak. . . Perhaps this is part of the reason we have progressed so little in our attempts at under­ standing the processes of double star genesis: we have concentrated our efforts on the exceptional freaks, and neglected the common man in space." nononfinnnonnnnnnnnnnnnnnnnnnnnonnnnnnonnnnnnnnfinno ♦NOTE: DATA IS ASSUMED TO BE LAMBDA VS FNU, AND CONTINUUM IS CONTINUUM AND FNU, VS LAMBDA TOBE ASSUMED IS DATA ♦NOTE: ♦H0DE=2: SIMPSON'S RULE INTEGRATION OF PARABOLICALLT—HAPPED, PARABOLICALLT—HAPPED, OF INTEGRATION POINTS. RULE DATA SIMPSON'S EQUAL-SPACED LINEARLY-HAPPED, OF ♦H0DE=2: INTEGRATION ♦ MOD^I: MD=: A CANLITGAIN T RPZIA PRX6 LINEAR 6 APPROX TRAPEZOIDAL BT INTEGRATION CHANNEL RAN ♦MODE=0: NRAV= i OF WAVELENGTH SETS TO BE AVAILABLE. UP TO SIX TO UP AVAILABLE. BE TO SETS WAVELENGTH OF i NRAV= WHPIF(NBPL.NE.0) P= O LT PCPLRSLST EHD. O ASSIGN TO HADE. BE TO RESULTS OPCHP'L PLOTS OF • NPL= NCDL= f COLUMNS USED IN OUTPUT TABLES OUTPUT IN USED COLUMNS f NCDL= RADIAL FOR CORRECTED BE CAN BANDPASSES IF(NVR.EQ.1) 1000) .LE. (ROST BE SPECTBA TAPE IN OFCHANNELS # NVR= NCHAN= T= O AE T EPOESD NT FNP1 AT NTP>11 IF NOTE PROCESSED. TO BE TAPES OF • ABOVE) (SEEDESCRIPTION INTEGRATION OF MANNER NTP= HODE= PARAMETERS: INPUT OPTIONAL "NDATE" OF VALUE TAPE ON GIVEN POINTS THE TO HAVING SYMBOLS SOLID ASSIGNS ST»45 "ISOLID" SEE 0-99: SYMBOL THE VARIABLE VERSATEC ASSIGNED BE CAN POINT DATA EACH #110 FOLLOWING STATEMENTS SEE PLOTS 6 UP TO FOR ,YAXES X ASSIGN TO NOTE THAT PLOT LABEL CARD IS SPLIT INTO TWO 90 COLUMN LINES LINES COLUMN 90 TWO INTO SPLIT IS CARD LABEL PLOT THAT NOTE 98 TO UP 11,BANGING WITH MUSTSTART UNITt'S TAPE THAT NOTE PLOTTED S KH/S COMBINED BE V-Hh'N.N RILL “NNNN= *NTP' TAPES FROM V INTEGERS: AS IN READ ABE VELOCITIES RADIAL NOTE THAT IF NTP>1 AT ST# ST*41 TO AT NTP>1 LOOPS 150 IF PROGRAM THAT NOTE SHIFTS VELOCITY WIDTHS RADIAL FOR EQUIVALENT OF HE CORRECTED CAN INSTEAD SUMS FLUX BANDPASSES GIVES THAT NOTERAVR OR W A R NEGATIVE THIS PROGRAM CALCULATES EQUIVALENT WIDTHS OR FLUX SUMS IN ONE OF ONE IN SUMS FLUX ORNATS: THREE WIDTHS EQUIVALENT CALCULATES PROGRAM THIS ♦ PORM IIO2 ♦ ♦♦♦ VIDICON2 PROGRAM ♦ ♦♦♦ DEFINED BT STRAIGHT LINES IN FNU VS NU OR LOG LOG OR NU VS (FNU) NU FNU VS IN LINES STRAIGHT BT DEFINED DATA. EQUAL-SPACED CONTINUUM. SEDSRPINO OFPRMTR BELOH) PARAMETER LOGF OF (SEEDESCRIPTION IE VLEO 'NDATE* OF VALUE GIVEN SPECTRA ON THE SAME TAPE. SEE 'IWAVE' AT ST«45 AT ST«45 'IWAVE' SEE TAPE. SAME THE ON LINE SPECTRA BROAD/NARROW FOR USEFUL PARTICULARLY IS OPTION EEAE, PT IE E PO. E »SR BPLOT' SEE PLOT. PER LINES 6 UPTO GENERATED, YBL09: E S*5 TE AIBE 'ISOLID* VARIABLE THE THE HAVING TAPE ON PTS TO SYMBOLS SOLID ST*45. ASSIGNS SEE 0-99: SYMBOL ST#110 FOLLOWING ST'S SEE PLOTS 6 UP TO FOR AXES X,Y B LOE (IH XRCINWNOSEC) THIS EACH). WINDOWS (WITH 6EXTRACTION ALLOWED ABE E SAE I S RLT "• (IN (ONCARD) RPLOT) SR PLOTPARM XSCALE SCALE SET LINES. EACH DATA POINT CAN BE ASSIGNED VERSATEC VERSATEC ASSIGNED BE CAN POINT DATA EACH COL'H 40 LINES. TWO INTO SPLIT IS CARD LABEL PLOT THAT NOTE MUST 98 TO UNIT#'S UP 11, RANGING TAPE WITH START PLOTTED. 0 BECOMBINED RILL NTP-10 TAPES FROM DATA ALL S ST#U1 TO LOOPS PROGRAM ST#150 IN READ AS ARE RV'S INTEGERS: SHIFTS. VELOCITY -NN VHNN HS SES# 101) (SEEST# KH/S ♦ V-HNN.N /-BNHN= Program Program VIDIC0N2 NBPL RESIDUAL INTENSITY PLOTS WILL BE WILL PLOTS INTENSITY RESIDUAL NBPL APPENDIX 120 t AL DATA ALL C C BEAD IN PLOT LABELS IP BEQ'D BEQ'D IP LABELS PLOT IN BEAD C oonoooo oooooo nooonooo 0 ED510 LBL(,)B11) (LABELB(H,N),B=1,10) BEAD(5,100)205 (LABELA(H,N),B-1,10), PORHAT(2F7.1,10) 551 200 CONTINUE 200 WAVR(H,N)=WAVR(H,N)/1.D1 201 5 PORBAT(12P6.0,210) 151 150 PORHAT(1110) 100PORHAT(20A0) NPTS(N)=0 199 01 IF(NT.NE.10)REWIND NT NT IF(NT.NE.10)REWIND 01 00 NT=10 00 IE 2CHAN,BODE 1ABELB(10,6) ,NHEAD(20) ,NDATE(10) ,NAME(10),NPILE,RRPL1 1ELX (10) ,LABELY (10) 1, 17) f IPIND (260) ,ISYH (260) ,IVRAD(260) ,IWAV(260) ,NCONT(6) ,BSTEP (6) ,N NT=NT+1 ZP(HODE.GE.O)GO TO 00TO ZP(HODE.GE.O)GO DO 205 N=1,BPL N=1,BPL 205 DO 01 TO IP(NPL.LE.O)GO IF(NPBDGE.EQ.I)READ(5,551)WCV,WCR,NHC IF(NPBDGE.EQ.I)READ(5,551)WCV,WCR,NHC BELOW LOOP 215" "DO SEE WCV,HCR= CONTAB WINDOW NNC^ICHANNELS TO AVE AT EACH CONT PT CONT EACH AT AVE TO NNC^ICHANNELS #1 WINDOW EXTRACTION TO APPLIED BE ONLY CAB WINDOW CONTAB WCV,HCR= HB: IF(NPDDG£.EQ.1.AND. IWAV(NP).LT.O). EQW BAIN PROMCALC»D BOW CONTAH OPP SOBTRACTS CONTAHINATION WING POR FIXOP N=1,6 201DO READ(5 ,151) (WAV B=1,NWAV 200 DO V (H,N) ,WAVB(B,N) ,H=1,6) ,NCONT(H) ,NSTEP (H) HAVV(H,N)=WAVV(B,N)/1.D1 HAVV(H,N)=WAVV(B,N)/1.D1 BEAD(5,150) NCHAN,BODE,NVR,NTP,NPL,NCOL,NWAV,NRPL,LOGP,NPDNCH,NPDDG READ(5,100) (NHEAD(N) ,N (LABEL1RBAD(5,100,EHD=90) = (N)1,20) ,H=1, 10) , (LABEL2 (N) ,H=1,10) DO 199 N=1,6 DO C=2.997925D5 S= • HSP=0 OPTIONS S HEADINGS TABLE OOTPDT 9 READ EGNRTD BWR: 1 PO S EEAE O EC SPECTRDH. EACH POR GENERATED (1)IS PLOT BEWARE: GENERATED. BE WILL PLOTS INTENSITY (HRPL.NE.O)RESIDDAL EQWIP S CALC'D *S ARE ISOLID/'JO DATA H•/ ISM DIHENSIOR (1000,6) ,NPTS(6) (1000,2,6)COBHON/TH/TABLE ,EW (9) ,WAVV(6,6) ,WAVB(6,6),LABELA (10,6),L COHHOH WAVCAL (1000),FLDX(1000),RCONV(1000),FCONV(1000),C, NDATA(260 WAVCAL (1000),FLDX(1000),RCONV(1000),FCONV(1000),C, COHHOH EDD O BNPS CARDS BANDPASS 6 TO DP READ COHBON/DATA/X (1000) rT (1000) tISTHBL(1000) ,LABEL1 (10) .LABEL2 (10) ,LAB HEAL*4 X,Y HEAL*4 I HP BEAL*8 LICIT (A-H,0-Z) NSTEP= NCONT= NPDDGE= NPDNCH= LOGP 2 WINDOW. IF(NSTEP.EQ.O) IT IS SET TO THE • OP PIXELS •OP THE TO SET IS IT IF(NSTEP.EQ.O) WINDOW. IP (SEESDH) SR (BODE EVEN HADE IS (STI200+).EQ.2) HCOHT COHTAB HING LIKE CORRECTS XF(NPODGE.EQ.I) GIVES PDHCHED CABD OOTPDT IP(9PDNCH.EQ.1): SEE ST112 SEE IP(9PDNCH.EQ.1): OOTPDT CABD PDHCHED GIVES LABBDA VS PHD OP DNITS IH BE TO ASSDHED SPECTBA TAPE • OP =SPACED STEPS TO BE IHTERPOL"D ACROSS EXTRACTED EXTRACTED ACROSS IHTERPOL"D BE TO STEPS =SPACED OP • AVERAGED BE TO POINTS CONTIHDDB ON CEHTERED PIXELS • IP(LOGP.EQ.I) COBTIHDOH HILL BE CALC*D HITH LOG LOG HITH (PHD)CALC*D BE HILL COBTIHDOH IP(LOGP.EQ.I) 121 non non noono nnnnn 0 PORHAT(10FB.1) 106 PORHAT(I3,1X, 105 10A3/10A3) 208 BRITE(6,103) (H,N,BAVV(H.N),BAVR(H,N),H=1,6) ,RCONT(H),NSTEP BRITE(6,103),RCONT(H),NSTEP 208 (H,N,BAVV(H.N),BAVR(H,N),H=1,6) (H) 0 PRA( OHROTOS NEFC: CA=,4» OE*I, BVR= HODE=*,I2,* NCHAN=»,I4,» EFFECT: IN OPTIONS OTHER PORHAT(• 104 103 PORHAT(• *,6(*B*,2I1, *(*,F7.1,•>*,F7.1,•)*,2X)/' CONT PTS AVE"D OV AVE"D PTS CONT *,6(*B*,2I1, *(*,F7.1,•>*,F7.1,•)*,2X)/' PORHAT(• 103 PORHAT 102 (* 1 (INANGSTBOHS)') *,20Att/*OINTEGRATION BINDOBS 207 IF(PCV.GT.O.DO.AND.PCV.LE.1.DO)PLOX(N)=PLDX (N)*PCV IF(PCV.GT.O.DO.AND.PCV.LE.1.DO)PLOX(N)=PLDX 207 NCP=NCP*1 206 0 POR101 HAT. II3,15 , 2I2,3A4, 914, 5A4) 152 PORHAT(4(F10.2,P10.5)) 152 55 READ(NT, READ(NT, 55 105,END=41) HPILE, (NDATE (R) ,H=1,10) , (NABE(N) ,N=1,10) 50 IP(NPR.LE.O)GO TO 52TO IP(NPR.LE.O)GO 50 45 READ(5 45 , 101, END=50) IFIND (NFR *1) ,IVBAD (NFR-H) ,ISTH (NFR+1) ,INAV (NFS*1 44 NFR=0 44 3 O27 N=1,NCHAN 207 DO 43 42 NC=NC+4 42 4A4,T92,2A4,T102,A4,T108,A4,T115,3A4) 22,* LOGP=*,12,* HPDHCH=',12,* NF0DGE=*,I2/*0FL«*,T7,‘DATE*,T14, NF0DGE=*,I2/*0FL«*,T7,‘DATE*,T14, HPDHCH=',12,* LOGP=*,12,* 22,* 3'TNAHE B S VRAD*,T31,3 (2A4,3X),T62,2A4,T70,A4,T75,2A4,T82,A4,T88, VRAD*,T31,3 (2A4,3X),T62,2A4,T70,A4,T75,2A4,T82,A4,T88, S B 3'TNAHE 1GE, (NHEAD(N) ,N=1,20) 1, HP ,2* P=,2* CL»I, NA=,2» NRPL=»,I NBAV=*,I2,» NCOL=»,I2,' 1 NPL=*,12,* * HTP= *,12,* ,12,* 1ER*,14,• CHANNELS. INTEGH HANGE DIV"D INTO',14,• INT"D STEPS*) INT"D INTO',14,• DIV"D HANGE INTEGH CHANNELS. 1ER*,14,• 1) ,(NDATA(NPR *1,N),N=1,17) BEAD TAPE DATA TAPE BEAD BEAD(NT,107) (PLOX (I),1=1,NCHAN) READ (NT, 106) (BAVCAL (N) ,N=1, BCHAH) BRITE(6,104)NCHAN,HODE,NVR,NTP,NPL,NCOL,NBAV,NRPL,LOGP,NPDNCH,NPDD BRITE(6,104)NCHAN,HODE,NVR,NTP,NPL,NCOL,NBAV,NRPL,LOGP,NPDNCH,NPDD H=1,NBAV 208DO IP(HODE.EQ.0)NSTEP (H)=0IP(HODE.EQ.0)NSTEP Rr(,0)(AE1(N)BRiro(6,102) (LABEL1 ,N=1, 10),(LABEL2(N),N=1,10) DPTTBE HEADINGS TABLE ODTPDT ISORT(NFR)CALL OT 45 TOGO NFR=NPR+1 IP(IBAV IP(IBAV (NFR+1) .GT.NBAV) IF(IHAV IBAV (NPR(NPR +IP(IPIND (NFR+ 1) 1) + .EQ.O) =1) IBAV (NPR+1) 1.LT.O) 50=1TO GO IP(IHAV(HPR IP(IHAV(HPR + 1*NBAV)1),LT.(— )IBAV(RFB + 1)=-1 READ(5,152) (BCONV (N) ,FCONV(H) ,N=HC,NC3) FLAG END OP PILE DECK BITH A FILEKO. A BITH DECK ARRAY) PILE (SEE OPNDATA END FLAG HERE ENTERED BE CAN DATA STELLAR ADD'L NOTE EDI PL EOS IT RDVLE, 6STHBOLi*S VALDES, VRAD LIST, BEQOEST PILE INBEAD NC3=NC+3 BCONV BITH SET (N)=0. DATA OP END PLAS IF IF (BCONV (N) H=NC,NC3 206 DO .EQ.O.DO) 43 TO GO PCV=PHAP2(BCONV PCV=PHAP2(BCONV rFCONV,NCP,B AVCAL) NC=-3 BCP=0 HAPPED IS SET DATA (IP BEQ DATA TO *D) BEAPPLIED TO SET CONVOLUTION IN BEAD IP(NT.NE.H) 44 TO GO IP(NT.NE.H) & CONVOLVED BITH SPECTRAL ARBAT POINT BT POINT BT POINT ARBAT SPECTRAL BITH CONVOLVED 122 O f DO 210 DO210 N=1,NFR NF=N NOTE NOTE THATCHANGES IN OUTPUT CALCS CAN BE BADE HERE GO TO GOTO 55 (NF)) V (IWA IW=IABS VRAD-0.D0 IF(NVR.EQ.O) IF(IPIND(H) .EQ.NFILE)GO TOIF(IPIND(H) 60 DON'T FORGET TO REASSIGN NCOL COLOHNS (• USED IN OUTPUT TABLE) (N).LT.O.DO)FLUX(N)=0.D0 IP(PLUX CHECK TAPE PILE* AGAINST REQUESTED LIST VRAD=DFLOAT (IVRAD (NF) )/1.D1 (NF) (IVRAD VRAD=DFLOAT VCORR=(VRAD/C)♦1.DO OUTPUT RESIDUAL INTENSITI PLOTS IF NRPL=1 SR SUB) (SEE IF(LOGF.EQ.I)GO TO 61 CORRECT INTEGR WINDOWS FOR VRAD CHANGE TO DLOGIO(FNU) UNITS IF(LOGP.EQ.1) DO 211 N=1,NCHAN GOTO 211 FLUX(N) =DLOG10(FLUX(N)) FLUX(N) NEQ=-1 NEQ=-1 DO215 N-1,6 NEQ=NEQ+2 IF(N.LE.3) WV=DABS(WAVV(IW,N))*VCOBB *VCOBB WR=DABS(WAVR(IW,N)) IF(N.LE.3)EW(NEQ)=EQW IF(N.LT.O.OR.EQW.LE.O.DO)GO TO 555 NRPL1=0 IF(WAVV(IW,N).EQ.O.DO)GO 215 TO NOEQW=0 IF(WAVV(IW,N).LT.0.D0.OR.WAVR(IW,N).LT.0.DO)NOEQW=1 NEQ=N+3 IF(N.GE.U) IF(N.LE.3)ER(NEQ+1)=DLAVE EH(NEQ)=-2.5D0*DLOG10(EQW) .EQ.ISOLID.AHD.IS.BQ.O)EW(NEQ)=EW(NEQ)-5.D0 (NDATE(4) IF SUBTRACT OFF CONTAB LINE(NFUDGE.EQ.1.AND.IWAV(NF).EQ. WING IN IF IF(NOEQW.EQ.0)NFPL1=NRPL1+1 CALL SUB(WV,WRfIW,N,NOEQW,LOGF,HRPL,EQWfDLAVE,NPR) TEBPORART FIXUP FOR JUN79 ND5 FLUX CORRECTION HAGNITUDES) (5 NNP-NRPL EQW WRITE(6,500) XF(IWAV(NF) .GT.O)GO TO 215XF(IWAV(NF) GO TO IF(IW.NE.1.0R.R.NE.1) 215 NNC1SNC0NT(IW) NCONT(IW)=NNC BRPL -EQW =EW(NEQ) EW(NEQ) CALL SUB(WCVfWCR(IWrN,NOEQW,LOGFrNBPL,EQN,DLAVE,NPR) 60 60 IS=ISTB(NF) 61 XP(PLUX (N) .LE.0.D0)FLUX(N) = 1.D-75 1.D-75 = .LE.0.D0)FLUX(N) (N) 61 XP(PLUX 107 107 FOBHAT(8E10.4) 210 210 CONTINUE 211 211 CONTINUE 555 IF(NFUDGE.NE.I)GOTO 215 500 FC»HAT{' SUBTRACTED CONTAB EQW FOB WINDOW #1:»,P5.1) uuu uuuuuuu u o u u non ooo oono 112 FOR112 FAT (I3,1X,2A3,1X,3A4,1X,9F6.2) 110 FOBHAT(• *,I3,lX,2(2A3r1X),I1,I2,P6.1,3(F6.2,F5.2),5F6.2,F6.3,IX,2 *,I3,lX,2(2A3r1X),I1,I2,P6.1,3(F6.2,F5.2),5F6.2,F6.3,IX,2 FOBHAT(• 110 CONTINUE 215 52 IF(NPL.LE.O) 52 89 TO GO IF(OHV.EQ.O.DO) 655 GO TO 5 IF(HG.EQ.O.DO)GO TO 6TO IF(HG.EQ.O.DO)GO 5 IF(DHV.EQ.O.DO) 45 GO TO 1,3),(EH(N) ,N 1,5,2)— (EH(N),N-7,NCOL) , ,BHAG,DBV,B6 1A4,1X,A4,1X,5A4) 1N),N=1,NC0L) ,BHAG,UHV,WG, (NDATA(HF,N),N=1,3), (BDATA(NF,N) ,N=13,17) TABLE(NPTS(6),2,6)=EH(3) TABLE(NPTS(6),1,6)=UHV O21 8=1,10 221DO N=1,NPL 220 DO OHTPOS FPORH ENDING IFPROGRAH PLOTS OOTHJT 55 TOGO ISH (NPTS (6) ,6)-ISP1 TABLE (NPTS (5) ,2 ,5) =EW (3) TABLE(NPTS(3) TABLE(NPTS(3) ,1 ,3)=BHAG TABLE(NPTS(2) ,1 ,2)=BHAG NPTS(2)=HPTS(2) +1 TABLE (NPTS (1) ,2 ,1) =EH (3) NPTS (6) =NPTS (6) +1 TABLE(NPTS(5),1,5)=BG ISH (NPTS (5)NPTS(5)=NPTS(5) *1 ,5)=ISP1 TABLE (NPTS (U) ,1,4)=UHV NPTS(4)=NPTS(4) +1 TABLE(NPTS(3),2,3)=EW(9) I S H ( N P T S ( 3NPTS(3)=NPTS (3)*1 ) , 3 ) TABLE(NPTS(2) =,2,2)=EB (7) I S P 1 9 TO IF(BHAG.EQ.O.DO)GO QUANTITIES ODTPDT OTHER PLOT TO STEPS THESE CHANGE TABLE(NPTS(4),2,4)=EW (8)TABLE(NPTS(4),2,4)=EW ISH (NPTS (4) ,4)=ISP1 ISH (NPTS (2) ,2)=ISP1 TABLE(NPTS(1) ,1,1)=EW(8) 55 TO IF(NPL.LE.O)GO HSP=NSP+1 1SH (NPTS (1) ,1)=ISP1 IF(NDATE(4).EQ.ISOLID)ISP1=-1*ISP1 ISP1=IS+1 IF(NPDNCP.EQ.1) DATA ODTPDT PUNCH NPTS(1)=NPTS(1) +1 IF(NPONCH.EQ.I) BRITE (7,112) NFILE, (NDATE BEQ'D IF (N) ARRAY LOADPLOT ,B=4,5) , (NDATA (NF,N) ,N=1 (NAHE (N) ,N=1,2),IW,IS,VRAD,(EW( WHITE(6,110)NFILE,(NDATE(N),N=4,5), WG=DFLOAT (NDATA (NF,9) J/1.D1 HRPL=NNP HRPL=NNP I I C O N T ( I B ) =NNC 1 DH7=DFLOAT(NDATA(NF,5)+HDATA(HF,6))/1.D2 DH7=DFLOAT(NDATA(NF,5)+HDATA(HF,6))/1.D2 BHAS=DFLOAT BHAS=DFLOAT (NDATA (HF,4) EW(8)=EB +NDATA(8) (NF,5) -EW J/1.D2 (9) 124 125

LABELX(B)=LABELA(H,R) 221 LABELY(B)=LABELB(fl,H) HR=HPTS (B) DO 222 H=1,NN XSYHBL (fl) =ISH (fl ,N) X (H)=TABLE(H,1,N) 222 Y (H) =TABLE (H , 2, H) 220 CALL VPLOT (NPTS (N) ) 89 WRITE(6,115) DSP 115 PORBAT('O',16,* SPECTRA PROCESSED') 90 STOP END 126

IEIHD, IEIHD, IVRAD, & ISTH ARRAYS INTO ASCENDING ORDER SOB ROUTINE ISORT(H) SOB IBPLICIT REAL*B (A-H,0-Z) IBPLICIT REAL*B ALSO SORTS RDATA ARRAY PERFORBS SHELL-BETZNER SORT TO ARRANGE B=N B=N REAL*** I,Y REAL*** COBBON WAVCAL(1000),FLOX(1000),BCONV(1000),ECONV(1000),C,NDATA(260 COBBON/DATA/X (1000),Y (1000) ,ISYBBL(1000),LABEL1(10),LABEL2(10),LAB (1000) COBBON/DATA/X(1000),Y RETORN IP(B.EQ.O) K=N-fl 1=0 J=1 J=1 IF (IPIND (I) .GT.IPIND (L) ) GO TO GOTO 310 ) (L) .GT.IPIND (I) (IPIND IF GO TO 100 IF(O.GT.K) I FIND (L) =IT =IT =IPIND(L) IFIND(I) (L) FIND I (I) IV=IVRAD IVBAD(L)=IV GO TO GOTO 200 IVRAD (I) =IVBAD(L) (I) IVRAD (I) IS=ISYB =IS (L) ISYB ISTB(I) =ISYB(L) ISTB(I) (L) V =IWA IB=IRAV(I) (I) V IWA IWAV(L)=IW DO 315 DO315 NN=1,17 (L,NN) =NDATA NN) (I, NDATA ID=NDATA(I,NN) GO TO 300 1) IP(I.LT. 1 = 1 -fl GOTO 250 END 1, 17),IPIND (260) ,ISYB (260) ,IVHAD (260) ,IWAV(260) ,BCONT(6) ,NSTEP (6) ,N ,N (6) ,NSTEP ,BCONT(6) ,IWAV(260) (260) ,IVHAD (260) ,ISYB (260) 17),IPIND 1, 1ELX (10) ,LABELY(10) (10) 1ELX 2CHAH,BODE

100 100 B=H/2 200 250 250 L=I+B 300 J=J+1 310 IT= IPIND (I) IPIND 310 IT= 315 NDATA(L,NN) =ID 315 NDATA(L,NN) uuuuu oonoo nnnnn onoo 200 CONTI 200 NOE 0 OHT* *,71*'F., OTEDT O ON: O SET=0***) EOB FOOND: NOT ENDPTS COHT **,,F7.1,*>',F7.1,' 100FOBHAT(* 0 PV=FLUX(RV)10 5 IP(NR«LE.NAVCAL(N-1)) IP(NR«LE.NAVCAL(N-1)) 200 TO GO 5 2CHAN.BODE 1ABELB (10,6) ,HHEAD(20) .NDATE (10)1ELX ,HABE (10)(10) .LABEL! ,NFILE,BKPL1 (10) 1, 17) .IPIND (260) .ISTH (260) .IVRAD (260) .IBAV (260) ,NCONT(6) .NSTEP (6) ,N N2=H1+NCONT N2=H1+NCONT (IB) -1 DLAH= (BAVCAL (NV+1) -BAVCAL (NV-1) )20 TO /2.D0 IF(NEHD.EQ.2)GO NEND=2>NR BEBD=1>NV, ENDPTS AT VALUE CONTINDUH AVE DETERBINE I1=NV-NCONT(IR)/2 I1=NV-NCONT(IR)/2 205NEHD=1,2 DO ENDPOINTS BEN ASSIGN 30 TO IP(NOEQB.EQ.1)GO IF(NCONT (IB) .LE.1) 30 TO GO IF(IABS(HODE).E0.2)NSTP=(NSTP/2)*2 NEWPTS=NSTP+1 IF(NSTEP (IB) ,LE.O)NSTP=NR-NV FR=FLUX(NR) RETURN OT 10 TO GO IF(HODE.EQ.O)NSTP=NR-NV IF(HODE.EQ.O)NSTP=NR-NV IF(BODE.EC.2) ROLE SIBPSOH'S FOR EVEN NSTEP BAKE FLDXES ENDPT SAVE BRITE(6,100)BV.BR NRH1=NP-1 IF( (BAVCAL(N) -CHHB) .GE.BR)NR=H-1 CHHB= (BAVCAL (N) -BAVCAL (N-1))/2.DO NOTE HODE=0 ADTOHATICALLT ASSIGNS NSTEP EVEN IF ALREADI SPECIFIED SPECIFIED ALREADI IF EVEN NSTEP ASSIGNS ADTOHATICALLT HODE=0 NOTE FB.TBVA()G O 200 TO IF(BR.GT.BAVCAL(R))GO IF((BAVCAL(N)+CHHB).LE,BV)NV=N+1 CHHW= (BAVCAL (N+1) -BAVCAL (N)) /2.D0 NR=N NV=N IF(NV .RE .0) 5 O T GO IF(BV 200.GE TO .BAVCAL IF(WV.LT.VAVCAL(N))GO (N + 1)) 200 TO GO SEARCH ENDPT VOF END FLAGS RV.RE.O REGION INTEGRATION OF CHANNELS ENDPT FIND E0BOO.DO COB BON/TN/T ABLE (1000,2,6) ,EB (9) W (6,6) .HA ,BAVR(6,6) .LABELA (10,6) ,L 0(1C HON BAVCAL(1000) .FLUX(IOOO) .BCONV (1000) .FCONV(IOOO) ,C,NDATA(260 O20 N=1,NCHAN 200DO NV=0 DLAVE=0.D0 NSTP=NSTEP(IB) EQW=0.DO COBHON/DATA/X (1000) ,Y (1000) .ISTHBL(1000) ,LABEL1 (10) .LABEL2 (10)X.T .LAB REALM REAL*B (A-H.O-Z)IHPLICIT SDH(BV.BR.IB,WB.HOEQB,LOGF,HRPL,EQW,DLAVE.HPR) SUBROUTINE 127 no non onnon non 210 CONTINOE 210 205 CONTINOE 205 206 IF(L0GF.NE.1)C0NT=C0NT+FL0X(N) 206 33 FRAT=FNOL/FNOC FRAT=FNOL/FNOC 33 32 DCND=C/HAYCAL(HR)-C/VAYCAL(NV) 32 30 DPN0RY=FL0X DPN0RY=FL0X 30 (NR) -FLOX (NY) 26 FLOX(NR)=CONT*DLAH/DLAHT 26 5 CONT=0.D0 25 20 H2=NR+NC0NT H2=NR+NC0NT 20 (IN) /2 EI HDs LOOP HODEs1 BEGIN OT 45 TO GO CALC AYE BESOLOTIOH AYE CALC Y (HPR) =FRAT DLA VE=DLAYE/ (NR-NY+1) HPR=N-NY + 1 X (HPR) =WAYCAL (N) 210 TO IF(NRPL.EQ.O)GO EQS=EQH+(1.D0-FRAT)*DLAN OT 210 TO GO EQS=EQW+FNOL*C*DLAH/NAYCAL(N) **2 EQS=EQW+FNOL*C*DLAH/NAYCAL(N) 33 TO IF(NOEQB.N£.1)GO CHO=C/SAYCAL (N) DLAVE=DLAVE+DLAH DLAVE=DLAVE+DLAH DLAS= (NAYCAL (M+1)IF(L0GF.EQ.1)FNUC=1.D1**FN0C -NAYCAL (N-1))/2.D0 3=FLOX(NR)C -CA*C/HAYCAL(NR) FNOC=CA*CNO*CB IF(LOGF.EQ.1)PNOL=1.D1**FNOL FNOL=FLOX(N) DO 210 N=NV,HR N=NV,HR 210 DO CA=DFHORY/DCNU (FNO=A*NO+B) EQ'N CONTINUDH LINEAR DETERKINE UNSPECIFIED IF ASSOHED=0 NODE IS BOTE TO GO .35 EI BD= LOOP B0DB=0 BEGIN DLAYE=(WR-WV)/NSTP PIB(OE.Q2G O 40 TO IP(IABS(HODE).EQ.2)GO SCALC®-1.D0*WAVCAL(NY)*DFNORV/DBAYBY XF(IABS (NODE)32 32 .GT.2)GO TO TO IF(BODE.EQ.O)GO DHAYR Y=HAYCAL (NR) -W A VCAL (NV) LIREAH COHTIHOOH SCALE FACTOR IH HO IH FACTOR SCALE COHTIHOOH LIREAH FLOX(NY)=CONT*DLAN/DLAHT OT 205 TOGO IF(LOGF.EQ.1)FLDX(NR)=DLOG10(FLOX(HE)) DLAH= (BAYCAL (HR*1) -WAYCAL (NB-1) )/2.00 IF(LOGF.EQ.1)FLOX(RY) IF(LOGF.EQ.1)FLOX(RY) =DLOG10(FLOX(HY) ) 26 TO IF(HEND.EQ.2)GO DLAHT=(NAYCAL (N2 + 1)-SAVCAL(H1-1)♦HATCAL(H2)-BAYCAL(HI))/2.DO IF(LOGF.EQ.1)COHT=CONT+1.D1**FLOX (H) IF(LOGF.EQ.1)COHT=CONT+1.D1**FLOX 206 N*N1,N2 DO 60 TO 25 TO60 H1=H2-NC0HT(IS)+1 129

C FIND FIRST WHEW BETWEEN NN C NN + 1 C LINEAR INTERPOLATION S INTEGRATION C LINEAR CONTINDOB IN LOGFNO VS NO IF(LOGF.EQ.1) C NOTE THAT IF(NCONT.GT.1)FRAT HAT HOT= 1.D0 AT INTERPOL'D ENDPTS C 35 NPH=HEWPTS DO 215 N=1, HEWPTS WHEW=WV+DFLOAT(N-1)*DLAVE NVfl 1=NV-1 DO 220 NN=NVH1,NR IF(WNEW.GE.WAVCAL(NN).AND.WHEW.LT.WAVCAL (NN + 1))GO TO 36 220 CONTINUE 36 SCALN0 = ( (W A VCAL (HN+1) -WHEW) *WAVCAL (NN))/( (WAVCAL(NN+1) -WAVCAL (NN) ) 1*WNEW) F NO INFLUX (NN + 1) ♦SCALNU* (FLOX (NN) -FLOX (NN + 1)) IF(LOGF.E0.1)FNOL=1.D1*+PNDL FNOO=FLOX(NR)♦ ( (WAVCAL(HR)-WNEW) *SCALC)/WNEW IF(LOGF.EQ.I)FNOC=1.D1**FNUC IF(NOEOW.NE.1)GO TO 39 EQW=EQW+FN0L*C/WNEW**2 GO TO 215 39 FRAT=FNOL/FNDC EQV=EQW+1.D0-FBAT IF(NRPL.EQ.O)GO TO 215 X (N)-WNEW T (N)=PRAT 215 CONTINOE EQW=EQW*DLAVE GO TO 45 C C BEGIN HODE=2 LOOP C PARABOLIC INTERPOLATION 6 SIHPSON'S ROLE INTEGRATION C LINEAR CONTINOOH IN LOGFNO VS NO IF(LOGF.EQ.1) C NOTE THAT IF(NCONT.GT.1)FRAT HAT NOT= 1.D0 AT INTERPOL’D ENDPTS C 40 NPR=NSTP+1 DO 225 N=2,NSTP,2 WNEW=WV+DFLOAT(N-1)*DLAVE WNEW1=WHEW+DLAVE IF(NOEQW.NE.1)GO TO 41 FN0C=WNEW**2/C FNDC1=WNEW1**2/C GO TO 42 41 FNOC=FLOX(NR)♦ (WAVCAL(NR)-WNEW)*SCALC/WNEW IF(LOGF.EO.I)FNOC=1.D1**FNOC FN0C1=FLOX(NR)+(WAVCAL(NR)-WNEW1)*SCALC/WN EW1 IF(LOGF.EO.1)FNOC1=1.D1**FNOC1 42 FN01FFHAP2(WAVCAL.FLOX,NCHAN,WNEW) IF(LOGF.RQ.1)FNOL=1.D1**FNDL FN0L1=FHAP2(WAVCAL,FLOX,NCHAN,WNEW1) IF (LOGF.EQ.1)FNOL1=1.D1**FNDL1 FRAT=FNOL/FNOC FRAT1=FNDL1/FHUC1 IF(HOEQW.EQ.1)GO TO 225 IF(N.RE.2)GO TO 43 FLV=FHAP2(WAVCAL,FLUX,NCH AN,WV) IF(LOGF.EQ.I)FLV=1.D1**FLV FCV=FLDX(HR)♦ (WAVCAL(HR)~WV) *SCALC/WV IF(LDGF.EQ.I)FCV*1.D1**FCV FRATV=FLV/FCV IF(NRPL.LE.O)GO TO 225 X (1)=BV T(1)=FRATV 43 IF(NBPL.LE.O) GO TO 225 X (N)=WNEW X (N+1)=BNEB1 T {N)=FRAT T (N+1) =FRAT1 225 EQW=EQW*4.D0*FRAT+2•D0*FRAT 1 IF (HOEQB.EQ.1)EQW=(EQB+(FLV/BV**2-FHUL1/WR**2)*C)*DLABE/3.D0 IF(N0EQB.HE.1)EQB=BR-BV-(BQB+FRATV-FRAT1)*DLAVE/3.D0 45 IF(NOEQB.HE.1.AHD.HRPL.GE.1)CALL RPLOT (NPR,IS,RB,BRPL) FUJX(NV) =FV FLOX(NR)=FR RETURN END non 50 FHAP2=A4 FHAP2=A4 50 (B+C*XNEB) *XNEH 0 IF(L.EQ.LL)30 50 TO GO 25 D=(FOLD D=(FOLD 25 (L) -FOLD (L1)) /(XOLD(L) -XOLD(LI) ) 22 C=CBAC C=CBAC 22 21 L2=L-2 21 0 IF(L.EQ.LL)20 50 TO GO 10 IF(XNEB.LT.XOLD 2010(L))GOIF(XNEB.LT.XOLD TO 2 (XOLD (L) -XOLD (1-1)) 1 ( F O L D ( L 1 ) / ( X O L D ( L * 1 ) -XOLD ( L 1 ) ) -FOLD ( L ) / ( X O L D (L+1 ) -XOLD ( L ) ) ) / (XOLD(LI)-XOLD(L2)) 2 1 (FOLD (L2) / (XOLD (L) D L -XO (L2)) D L O F - (L1)/ (XOLD (L) D L -XO (LI)))/ LL*=L OT 50 TO GO LL=L C=CPOB*WT*(CBAC-CFOR) IF (DABS (CFOB) .BE.0.D0) NT=DABS (CFOB) /(DABS (CFOB) BF0R=D- (XOLD4-DABS (L) +X0LD (CBAC)(LI)) ) *CFOB END RET tJEN A “FOLD (L) = B (FOLD -XOLD(L) (L)*B C=0.D0 D L O -F (L-1) ) / (XOLD (L) D L O X - (L-1)) B=8F0R+BT*(BBAC-BFOR) NT“0.D0 L = f l I N 0 (NOLD A=AFOR*BT*(ABAC-AFOR) CFOR=FOLD(L*1)/ (CFOR=FOLD(L*1)/ (XOLD(L+1)-X0LD(L)) (XOLD(L+1)-XOLD(L1))) ♦ * 50 TOGO LI^L AFDB=FOLD AFDB=FOLD (L1) -XOLD (LI) *D*XOLD (L) *XOLD (LI) *CFOB A“ABAC A“ABAC B=BBAC D L O F = C A B C (L) / ((XOLD (L) -XOLD (LI) 25 TO GO ) * (XOLD (L) D L O -X (L2))) + IF(L.LT.NOLD)GO TO 25 TO IF(L.LT.NOLD)GO ABAC=FOLD(L2) -XOLD (L2)(XOLD BBAOD- (L1) *D*XOLD (L1) 4XOLD (L2) *XOLD (L2) ) *CBAC *CBAC B= B= ( F O L D ( L I ) -FOLD ( L 2 ) ) / ( X O L D ( L I ) -XOLD ( L 2 ) ) CBAC=CFOR IF(L.EQ.NOLD)GO TO 22 TO IF(L.EQ.NOLD)GO ABAC=APOR BBAO=BFOR 10 TOGO LL=0 IF(L.GT.LL*1.0R.L.EQ.3)GO TO 21 TO IF(L.GT.LL*1.0R.L.EQ.3)GO LT=L-1 ZF(L.SQ.2) 30 TO GO 30 TO IF(L.GT.NOLD)GO L=L*1 L=2 BEAL*** X,T DIHEN SIONXOLD(1) ,FOLD (1) REAL*B (A-H.O-Z)XHPLICIT KDBOCZ)PBOGRAfl BT (SEE ATLAS6 INTERPOLATION PARABOLIC OCIH FHAP2{XOLD,FOLD,*OLD,XNEB) FONCTIOH , L) 131 NP1FA FRFRTRS INTPLOT RES FIRST FOR NRPL1=FLAG GENERATED BE NRPIr= TO PLOTS INT RES OF NUHBER TOTAL C C TAPE PROH SPECTROH EACH FOR CALLED GENERATED IS PLOT A THAT ROTE C H= UBR FETATDLN IDW (NN.LE.6) WINDOW LINE EXTRACTED OF (IW.LE.6) NUHBER SET C WINDOW EXTRACTION OF HR= NUHBER PLOTTED BE TO IW=LINE IB POINTS DATA OF NUHBER NPB= OVERLAP AVOID TO C PROGBAH HAIN IN FIRST LINES INPUT BROADEST C HODE PLOTCV IN C WORK NOT WILL ROUTINE PLOTV THIS C C C C onnnn IN CALC'D EQW*S IE HAIN PLOTS IBTENSITI RESIDUAL PROGRAM GBJERATBS C 200 THtN=AHIN1 THtN=AHIN1 200 (7HIN,T (N) ) XHAX=X 6 (1) 5 TZERO=TZERO+TSTEP 5 1ABELB(10,6) ,HHEAD(20).NDATE(10) ,HAHE(10),NFILE,NRPL1 1ELX (10) rLABEL! (10) CALL STHBOL(9.,.25,.07,LAHEL2,0.,+90) CALL CALL STHBOL(9.,.39,.07.LABEL1,0.,*90) CALL STHBOL(.25,.25,.21,NAHE1,0.,+3) CALL STHBOL(.25,.50,«19,NDATE9,0.,+3) CALL CALL STHBOL(.88,.25,.21,NAHE2,0.,+3) STHBOL(.67,.50,.19,HDATE5,0.,+3) CALL NAHE2=NAHE(2) STHBOL(2.99,-.25,.19,»ANGSTROHS•,0.,*9) CALL CALL AXIS(0.,0.,'(RINT-RCORE) AXIS(0.,0.,'(RINT-RCORE) /(1-BCOBE) CALL •,*22,9.,90.,0.,.25) NAHE1=NAHE(1) NDATE5=RDATE (5) NDATE9=NDATE(9) NUHBER(.67,.68,.14.FIOAT(NFILE),0.,-1) CALL STHBOL(.25,.68,.19,»FL«•,0.,+3) CALL PLOT(7.,0.,+2) CALL *,-1,4.,90.,0.,.25) AXIS(7.,0.,' CALL NERPEN(3) CALL CALL PLOT(0.,0.,+3) CALL 10 TO IF(NRPL1.NE.1)GO WSCALE=(XHAX-XBIN)/(7.*2.59*XSCALE) XHIN=AHIN1 (XBIN.X (N) ) I!CN=I (1) X STEP=0 • ALAI(.TEO• • AXIS(0.,TZERO,• CALL XHA X=AHAX1(XHAX,X(N)) OT 6 TOGO TZERO=9. IP(NRP.L1.NE.1)GO 5 TO DO 200 N=2,NPR N=2,NPR 200DO (1)TKN=T IF(NRPL.NE.1) TS TEP= 5./FLOAT(N BP L-1) PLOTS(0,0,0) CALL XSCALE=.857193 CARD! PLOTPARH ON SCALE SETXSCALE= AXES ANNOTATE 6DRAB COHHON/TN/TABLE COHHON/TN/TABLE (1000,2,6),EW(9) ,WAVV(6,6),RAVR (6,6),LABELA(10,6) ,L COHHON/DATA/X (1000) (10),LABEL2(10),LAB ,T(1000),ISTHBL(1000),LABEL1 REAL*9 IHPLICIT (A-H.O-Z) SET VERTICAL AXIS POSITION POSITION AXIS SET VERTICAL TABLE,ER.WAVV.WAVR REAL*8 SUBROUTINE RPLOT(NPR,IB,NW,NRPL) SUBROUTINE ,*1 ,7.,0.,XHIN,(XHAX-XHIN)/7.) ,7.,0.,XHIN,(XHAX-XHIN)/7.) 132 GO TO 15 10 CALL PLOT(0•, TZERO,*3) CALL PLOT(0.»TZERO-T5TEP,*2) CALL PLOT(7.,TZERO-TSTEP,*3) CALL PLOT(7.,TZERO,*2) 15 CALL STHBOL (5.5,TZERO-.5, .07, *R*,0.,+1) CALL HOHBEB(5.57,T Z E R O - . 5 07,FLOAT(IW),0.,-1) CALL HOHBER(5.6«,TZERO-.5,.07,PLOAT(NW),0.,-1) CALL STHBOL(5.5,TZBRO-.6«,.07,•BCOBE=*,0.,+6) CALL HOHBEB(5.92,TZERO-.6U,.07,THIN,0.,+3) CALL STHBOL(5.5,TZERO-.78,.07,•A/CH=*,0.,♦S) CALL HOHBER(5.85,TZERO-.78,.07,BSCALE,0.,+3) CALL NERPEN (1) CALL PLOT (0 . ,TZEBO, + 3) DO 205 R=1,NPR XPL=((X(H)-THIN) /(XHAX—XHIR)) *7. TPL=TZERO-((1.-T (N))/(1.-THIN))*U. 205 CALL STHBOL(XPL,TPL,.07,NH,0.,-2) IF(HRPL1.EQ.HRPL)CALL PLOT{0.,0.,+999) RETORN END onno non nonn 0 CONTINUE 200 52 CALL STBBOL STBBOL CALL 52 (0.0,8.0,0. 1U,LABEL1,0.,I»0) 51 CALL PLOT(X0+.07,T0,42) CALL 51 50 CALL PLOT(XO,T04.07,+2) CALL 50 1ELX (10) ,LABELT(10) END RETURN PLOT CALL (0.,0. ,+999) STHBOL(XPLOT,TPLOT,.19,NS,0.,-1) CALL IF(ISTHBL(N) STHBOL(XPLOT,TPLOT,.1U,3,0.,-1) .LT.O)CALL TPLOT= (T (N) XPLOT= (X -T (NPTS+1)(N) ) -X (NPTS/I (NPTS+2)+ 1) )/X (NPTS *2) NS=IABS(ISTHBL(N))-1 N=1,NPTS 200 DO RETURN S DATA THE PLOT .,Q0) STHBOL(0.0,7.7,0.19,LABEL2,0 CALL 51 TO GO IF ISTHBL(N).LT.O STHBOLS HILL BE FILLED IN BT OTERPLOTTING OTERPLOTTING BT IN FILLED BE HILL STHBOLS ISTHBL(N).LT.O IF PLOT(XO,TO,+2) CALL T0=T0-1. IF(TO.EQ.O.)GO TO 52TO IF(TO.EQ.O.)GO PLOT(XO,TO,*3) CALL 50 TOGO PLOT(XO,TO,+2) CALL PLOT(XO,TO,+3) CALL XO=X041. IF(X0.EQ.7.) 51 TO GO CALL PLOT(10,TO,+3) CALL 10=7. AXIS(0.,0.,LABELT,«-90,7.,90.,T CALL (NPTS + AXIS(0.,0. CALL ,LABBLX,-**0,7.,0.,X(NPTS1),T(NPTS+2) ) + HEWPEN CALL (3)1) ,X(NPTS*2)) SCALE(X,7,,HPTS,+1) CALL X0=0. AXES PLOT ANNOTATE &DRAH SCALE(T,7.,NPTS,+1) CALL PLOTS(0,0,0) CALL THIS PLOTV ROUTINE HILL ROT WORK IN PLOTCT RODE RODE PLOTCT IN ROT WORK HILL ROUTINE PLOTV THIS PL CT (A**H,0~Z) I HP LICIT COHBON/DATA/X COHBON/DATA/X (1000) ,T(1000) ,ISTBBL(1000) ,LABEL1 (10) ,LABEL2(10) ,LAB INITIALIZE FOR PLOTTING, DETERMINE SCALING FACTORS SCALING DETERMINE PLOTTING, FOR INITIALIZE SUBROUTINE TPLOT TPLOT (NPTS)SUBROUTINE 134 LIST OF REFERENCES

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