The Balmer decrement in the emission spectra of astronomical objects

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Authors Bloom, Gary Stuart, 1940-

Publisher The University of Arizona.

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Link to Item http://hdl.handle.net/10150/318347 THE BALMER DECREMENT IN THE EMISSION SPECTRA OF ASTRONOMICAL OBJECTS,

by Gary Stuart Bloom

A Thesis Submitted to the Faculty of the DEPARTMENT OF PHYSICS In Partial Fulfillment of the Requirements For the Degree of MASTER OF SCIENCE In the Graduate College THE UNIVERSITY OF ARIZONA

1 9 6 9 STATEMENT BY AUTHOR

This thesis has been submitted in partial ful­ fillment of requirements for an advanced degree at The University of Arizona and is deposited in the University Library to be made available to borrowers under rules of the Library. Brief quotations from this thesis are allowable without special permission, provided that accurate acknowl­ edgment of source is made. Requests for permission for extended quotation from or reproduction of this manuscript in whole or in part may be granted by the head of the major department or the Dean of the Graduate College when in his judgment the proposed use of the material is in the interests of scholarship. In all other instances, however, permission must be obtained from the author.

SIGNED:__

APPROVAL BY THESIS DIRECTOR This thesis has been approved on the date shown below: PREFACE

Due to the nature of the problem undertaken, namely, comparison of the Balmer decrements In astronomical emission spectra, the details of this thesis were becoming dated even prior to completion of the work. We have tried to gather together results of all studies pertinent to this report which have been published in English, German, or French prior to the fall of 1 9 6 5. Some studies published in Russian are included, but we do not claim to have exhausted this literature. Work published in other languages have not been included. Due to the necessity of completing this study, we have ignored the promulgation of results currently being published. Inclusion of such data should give more accurate results for some of the objects we examined, and would probably improve the averages which we have computed. At least one such study, Vorontsov-Vel' yaminov et a l , (1965), gives corrections of their former data which we have used in this study. Their recalibration gave them intensity increases upward of 20$. In our search into the literature of objects exhibit­ ing Balmer emission spectra, we inadvertantly overlooked the Wolf-Rayet class of stars which are known to exhibit these

ill lines (see5 for example., Ambartsumyan 1958, pp. 490-496). A cursory search to remedy our lack of information about the decrement has not been successful, and we have found only an emission decrement from the nebulosity associated with a WC star. (Andrillat and Andrillat 1953)° TABLE OF CONTENTS PAGE LIST OF TABLES . viii LIST OF ILLUSTRATIONS xvil

ABSTRACT » . . » o 0 0 0 Q . . XX

INTRODUCTION . , 1 GENERAL REMARKS 4 The Definition ...... 4 Problems of Measurement ...... , 4 Corrections to the Raw Data . . . . , 6 The Uses of the Balmer Decrement 10 Types of Intensity Data Found in the Literature . . , 11 Normalizing the Data 14 GASEOUS NEBULAE...... 15 Historical Remarks 15 The Published Data 17 NGC 1535 .... 20 NGC 1976 . . , 22 NGC 2022 .... 22 ■ NGC 2392 .... 27 NGC 2440 ...... 28 NGC 6543 ...... 29 NGC 6572 .... 29 NGC 6741 .... 29 NGC 6826 .... 31 NGC 7009 .... 32 NGC 702? .... NGC 7662 .... 1 IC 351...... 36 IC 418 .....

10 2l49 o 9 a 00 . o © V e e » 38

IC 4997 o6 . a . o 0 *0. . » Implications of the :Published Data. 1: Average Decrements for Each Nebula 40 Steep, Average, and Shallow Decrements 45 Summary and Conclusions ...... 55

v vi TABLE OP CONTENTS--Continued . PAGE STARS f 58 Observational Difficulties and Star Types . . . . 58 Be Stars 61 Variable Stars 62 Pulsating Variablea--.3 Canis Major is Type . , 66 Eruptive Variables-=T Tauri Stars ...... 67 Eruptive Variables--Novae ...... i 67 Description ...... 67 Published Data ...... 71 Eruptive Variables “--Nova-Like Variables . . . 78 Recurrent Novae ...... 81 R Coronae Borealis Stars ...... 81 P Cygni Stars ...... 84 Symbiotic Stars ...... 92 Summary and Evaluation of the Published Data . . . 106 Additional Data ...... 110 SOLAR PHENOMENA ...... \ . . . . 112 Solar Regions and the Sites of Hydrogen Emission Spectra ...... 112 The Flash Spectrum of the Chromosphere ...... 115 History and Observational Techniques .... 115 Specification of Balmer Decrements in the Chromosphere ...... 122 The Published Data ...... 128 Normal Regions ...... 128 Excited Regions ...... 133 Discussion of the Data and Some Comparisons ...... 134 The Average Chromospheric • Balmer Decrement ...... 163 Local Phenomena . . , ...... 186 Definitions ...... 186 Presentation of Data and Comments ...... 189 ' Quiescent Prominences ...... 191 Active Prominences ...... 193 Published Data ...... 193 Analysis of the Data ...... 196 Limb Flares ...... 200 Dxsk Flares . . . . . ■. @ ...... 201 Summary ...... ® 207 vii TABLE OF COEFEIfPS - -Continued PAGE COMCLTOIMG REMABl^ ...... 210 Galaxies ...... 210 Quasi-Stellar Objects ...... 212 tE> y ...... 212 A Suggestion ...... 213 APPMBXX A— Balmer Line Intensity Data as Originally Published for Nebulae ...... 215 APPENDIX B— Balmer Line Intensity Data as Originally Published for Stars ...... 219 APPENDIX C— Balmer Line Intensity Data as Originally Published for the Solar Chromosphere ...... 233 APPENDIX D— Intermediate Steps in the Computer Calculations ...... 251 APPENDIX E— Balmer Line Intensity Data as Originally Published for Solar Flares and Prominences .... 259 BIBLIOGRAPHY ...... 266 LIST OF TABLES TABLE PAGE 1. Nebulae Examined in Detail with References Thereto ...... 19 2. Balmer Decrements Published for NGG 1535 ..... 21 3. Averages of Mendez1 11 Point Observations of NGG 1976 23 4. Balmer Decrements Published for NGG 1976 24 5. Balmer Decrements Published for HOC 2022 ..... 27 6. Balmer Decrements Published for NGG 2392 ..... 27 7. Balmer Decrements Published for NGG 2440 ..... 28 8. Balmer Decrements Published for NGG 6543 ..... 29 9. Balmer Decrements Published for NGG 6572*...... 30 10. Balmer Decrements Published for NGG 6741 ..... 31 11. Balmer Decrements Published for NGG 6826 ..... 32 12. Balmer Decrements Published for NGG 7009 ...... 33 i 13. Balmer Decrements Published for NGG 7027 ..... 34 14. Balmer Decrements Published for NGG 7662 ..... 35.

15. Balmer Decrements Published for IC 351 ...... 36 16. Balmer Decrements Published for 10 4l8 ...... 37

17. Balmer Decrements Published for IG 2149 ..... 38 18. Balmer Decrements Published for 10 4997 ..... 39 19. Weights Assigned to Publications in Order to Compute the Average Balmer Decrement for Each Nebula ...... 44 viii ix LIST OF TABLES— Continued ' TABLE PAGE 20. Weighted Averages of the Balraer Intensities for Selected Nebulae {First Group) ...... 46 21. Weighted Averages of the Balraer Intensities for Selected Nebulae (Second Group) ...... 47 22. Average Balraer Line Decrements for Gaseous Nebulae ...... 50 23. The "Average11 Intensity Range in Gaseous Nebulae for Each Balmer Line and the Number of Nebulae Whose Intensities Fell Within and to Either Side of this Range ...» 52 24. A Comparison of Nebula Balmer-Line Intensities with the Group of "Average".Intensities .... 53

2 5 . Be Stars Examined and References Thereto ..... 61 26. Balmer Line Intensities and Decrements for y Cassiopeiae from 9/37 to 12/38 ...... 63 27. Balmer Decrements of 10 Be Stars ...... 64

28. Balmer Decrement of S Cephei ...... 67

2 9. Balmer Decrements of Two T Taurl Stars ...... 6 9. 30. Studies Made of Novae, Published Since. 1950 . , . 72 31. Balmer Decrements of Nova Serp 1909 and Nova ' Scu 1949 ...... 0 . .75 3 2 . Balmer Decrements of Nova Aqu 1918 » ...... 76 33. Balmer Decrements of Nova DK Lae 1950 ...... 77

34. Balmer Decrements of Nova Herculis i960 ..... 78 35. Balmer Decrements of Nova Herculis 1934 ..... 79 3 6. Balmer Decrements of Nova T Coronae Borealis . . . 82 37. Balmer Decrements of Nova RS Ophluchi 1958 .... 82

3 8 . Balmer Decrements of xx Ophiuchi (HD 116114) . . . 85 •X LIST OF TABLES--Continued TABLE PAGE

39. Balmer Decrements of P Cyg and HD 51585 ...... 89 40. Balmer Decrements of x Ophiuchl ...... 90 41. Studies Made of Symbiotic Stars Published Since 1950 ...... 93 42. Balmer Decrements of Z Andromedae ...... 97

4 3 . Balmer Decrements of BF Cygni ...... 99 44. Balmer Decrements of WC 6 0 3 ...... 100

4 5 . Balmer Decrements of 01 Cygni ...... 100 46. Balmer Decrements of x C y g n i ...... 101 4 7 . Balmer Decrements of AXPe r s e i ...... 101 48. Balmer Decrements of AG Pegasi ...... 102

4 9 . Balmer Decrements of BDtll°4673 ...... 103 50. Balmer Emission Line Intensity Value Spreads Published for Eight Star. Categories ...... 107 51. Studies of Stars Published Prior to 1950 , . . . . Ill 52. Balmer Line Studies Resulting from Solar Eclipse Observations ...... 117 53- Balmer Decrements Published for the Solar Chromosphere without Limb Height Determination ...... 135 54. Balmer Decrements Published for the 1932 Eclipse by Cillid and Menzel ...... 136 55. Balmer Decrements Published for the 1936 Eclipse by Athay, Menzel, and Orrall ...... 137

5 6. Balmer Decrements Published for the 1941 Eclipse by Vyazanitsyn...... 138 57* Balmer Decrements Published for the 1945 Eclipse by Vyazanitsyn ...... 143 z xi List of Tables---Continued TABLE PAGE

5 8 . Balmer Decrements Published for the 1952 Eclipse by Hontgast ...... 1 # 59• Balmer Decrements Published for the 1954 Eclipse by Houtgast ...... 144 60. Balmer Decrements Published for the 1952 Eclipse by Athay5 Billings* Evens and Roberts ...... 145 61. Balmer Decrements Published for the 1954 Eclipse by Martynov and Abluseva ...... 147 62. Balmer Decrements Published for Excited Regions During Eclipse ...... 148

6 3. Weights Assigned to Publications in Order to Compute lv (h0 ) and ln (ia0 ) ...... 166 64. Average Weighted Logarithmic Decrements for Three Representative Limb Heights ...... 167 6 5. Average Weighted Decay Coefficients ^'(hj) for Three Representative Limb Heights » . . . . 168 66. Average Logarithmic Balmer Decrements Computed for 14 Limb H e i g h t s ...... 170

6 7. Balmer Decrements at Selected Limb Heights Extrapolated from 1^(1000) and B^'CloO'^ .... 178 68. Balmer Decrements at Selected Limb Heights Extrapolated from In (2000) and B^JSOOO) .... 179

6 9. Balmer Decrements at Selected Limb Heights Extrapolated from ln ( 4000) and .... 180

7 0 . The Average Balmer Decrements at Selected Heights , ...... 181 71. Decrements from Early Observations of Local Phenomena ...... 190 72. Ratio of Central Intensities to Equivalent Widths in Quiescent Prominences as Reported by Ellison and Reid . . , ...... 191

I . xii BlSf OF CABLES — Oontlnued TABLE PAGE 73• Published Balmer Decrements for Quiescent Prominences ...... 194 74. Published Balmer Decrements for Active Prominences ...... 197 75- A Limited Average Balmer Decrement for Active Prominences ...... 199

7 6. Slopes of the Logarithmic Decrements in Two Active Prominences ...... 199 77. Published Balmer Decrements for Limb Flares . . . 202

7 8 . Published Balmer Decrements for Flares Seen Against the Disk ...... 205 79• Balmer Emission Line Intensity Spreads Published for Local Solar Phenomena ...... 209 80. Published Balmer Decrements for Seyfert Galaxies ...... 211 81. A Published Balmer Decrement for a QSO . . . « . . 212

82. Balmer Line Intensities for Nebulae5 Published by Kaler (1964) ...... 216

8 3 . Balmer Line Intensity Data for Nebulae# Published by O'Dell'(1 9 62) ...... 218 „ 84, Balmer Line Intensities for J320# Published by Burgess (1958) ...... 218

8 5 . Balmer Line Intensities for y Gass# Published by Wellman (1952) ...... 220 86. Balmer Line Intensities for Be Stars» Published by Burbidge and Burbidge (1953a#b) and Herbig (i9 6 0) ...... 221

8 7 . Balmer Line Intensities for 0 Oephei# Published by Wilson and Seddon (1956) ...... 222 88. Balmer Line Intensities for Nova RT Serp 1909# Published by Grandjean (1952) ...... 222 LIST OF TABLES--Continued TABLE PAGE

8 9. Balmer Line Intensities for Nova Sen 1949, Published by Colacevich (1950) ...... 222 90. Balmer Line Intensities for McRae +43°1, Published by Greenstein (1954) ... 222 91. Balmer Line Intensities for Nova Aqu 1918, Published.by McLaughlin (1953) ...... 223 92. Balmer Line Intensities for Nova DQ Herculis 1934, Published by Swings and Jose (1952) and Bajenov (1056) ...... 223 93• Balmer Line Intensities for DK Lacertae 1950, Published by Wellman (1951) 224 94. Balmer Line Intensities for RS Oph 1958, Published by Griffin and Thackery (1958) . . . .225 95. ' Balmer Line Intensities for T Cor Bor, Published by McLaughlin (1953) . 225

9 9 6. Balmer Line Intensities Published for Two P Cygni Stars ...... 226 97. Balmer Line Intensities for x Oph Published by Burbidge and Burbidge (1955) and Kupo (1959) 227 9 8. Balmer Line Intensities for BF Cyg Published by Merrill (1950) and Aller ( 1 9 5 5 ) ...... 228 99• Balmer Line Intensities for Cl Cygni Published by Merrill (1950) and Aller (1955) ...... 229 100. Balmer Line Intensities for x Cygni Published by Mao-Lin (1950b) and Fugita (1954) ...... 230 101. Balmer Line Intensities for AG Pegasi Published by Mao-Lin (1950a), Mao-Lin and Bloch (1952), Burbidge and Burbidge (1954), and Arkhipova and Dokuchaeva (1962) 231

102. Balmer Line Intensities for MWC 603 Published' by Tifft and Greenstein (1958) ...... 232 xiv LIST OF TABLES--Continued TABLE PAGE 103. Balmer Line Intensities for the Solar Chromosphere as Given In Some Early Publications ...... 234 104. Logarithms of the Intensities of the Balmer Lines in the Solar Chromosphere as Published in Cilli<§ and Menzel (1936) ...... 235 105. Logarithms of the Intensities of the Balmer Lines in an Excited Region of the Solar Chromosphere as Published In Cillid and Menzel (1936) ...... 236 106. Logarithms of the Intensities of the Balmer Lines in the Solar Chromosphere as Published in Yyazanitsyn (1951) ...... 237 107. Logarithms of the Intensities of the Balmer Lines in the Solar Chromosphere as Published in Vyazanitsyn (1952) ...... 242 108. Logarithms of the Intensities of the Balmer Lines in the Solar Chromosphere as Published ' in Athay, Menzel, and Orrall (1957) ...... 243 109. Logarithms of the Intensities of the Balmer Lines in the Solar Chromosphere as Published in Houtgast (1957) ...... 244 no. Logarithms of the Intensities of the Balmer Lines in the Solar Chromosphere as Published in Houtgast (1 9 62) ...... 244 111. Logarithms of the Intensities of the Balmer Lines in the Solar Chromosphere as Published In Thomas and Athay (1961) ...... 245 m . Logarithms of the Intensities of the Balmer Lines in an Excited Region of the Solar Chromosphere as Published in Thoms and Athay (1961) ...... 248

113. Logarithms of the Intensities of the Balmer Lines in the Solar Chromosphere as Published . in Martynov and Abluseva (1962) ...... 249 XV , LIST OF TABLES— Gont inued TABLE ' PAGE IIA. The Data inserted into the Computer Program . . . 252 115. The Fortran Program Used to Obtain Tables of Jlog^nlnXh)} from One Value of. &n(h0 ) and InTh©) for Each n ...... 253 116. The Fortran Program Used to Obtain Tables of fIn (h), 1 f' r7'KYT from One Value of % ( % ) and In ( h0 ) for Each n ...... 253

1 1 7. The Fortran Program Used to Obtain the Average Clog10In (h)3 ...... 254 IIB. The Fortran Program Used to Obtain the Average

{ t i f t } • • • • • • • • • • ...... • 255 1 1 9. Values of lo-fs. Computed from 0L(l6odkm) and In (i000km) 256 120. Values o f l o g 2knIn (h) Computed from Pn (20d0km) and InTsoSW) ...... 257 121. Values of logittIn(h) Computed from 4odokm') and I^SSoH) ...... 258 122. Balmer Line Intensities for Prominences as Published by Sirin and Tandberg-lanssen (I960) . o ...... 260 123-. Balmer Line Intensities for a Prominence as Published by Elliott, Ellison, and Reid (I960) . o ...... 260 124, Balmer Line Intensities for Proiiiinences~as Published by Shih-Huei (1961) ...... 26l 12:5. Logarithms of the Intensities of Balmer Lines . in Prominences as Published by UnsUld (1947) and by Athay and Orrall (1957) 262 126. intensities of the Balmer Lines for Solar Flares v Seen Against the Disk as Published in ■, Smith (1963) ...... 263

■t LIST OP TABLES— Continued TABLE PAGE 127. Logarithms of the Intensities of the Balmer Lines for a Solar Flare Seen Against the Disk as Published in Svestka (i9 6 0) ...... 265 LIST OF ILLUSTRATIONS FIGURE PAGE

1. A Schematic Reddening Curve ...... 9 2. The Three Common Measures of Emission Line Intensity ...... 12

3 .. Schematic Representation of Shallow* Average and Steep Balmer Decrements ...... 48 4. Nebular Balmer Line Intensities, as a Function of n ...... 51 5. Logarithmic Plot of Average Nebular Balmer Line Intensities as a Function of n ...... 57

6.. Balmer Decrements for Some Be Stars ...... 65 7. The Balmer Decrement of a B Canis Majoris Star . . 68 8. ^Imer Decrements of Two T Tauri Stars ...... 70

9. Typical Balmer Decrements for Some Novae . . . . . 80 10. Representative Balmer Decrements for Two Recurrent Novae ...... 83 11. Balmer Decrements of an R Coronas Borealis Star . 86 12. Characteristics of P Cygnl Stars11 Emission Line Profiles ...... e. 87

1 3 . Balmer Decrements of P Cyg Stars ...... 91 14. Representative Balmer Decrements of Some Symbiotic Stars (A) . 104

1 5. Representative Balmer Decrements of Some Symbiotic Stars (B) ...... 105 16. Schematic Representations of the Contacts of Lunar and Solar Limbs dufing Total Solar Eclipse ...... , 0 ...... 114 xvii - xviii LIST OF ZLLPSTRAT XOHS--Cont inued FIGURE PAGE 17. Schematic Drawing Showing Spectral Resolution of the Solar Image During Eclipse Totality . . . 119 18. Some Representative Logarithmic Decrements of the Chromosphere Seen in the 1941 Eclipse by Vyazanitsyn ...... 149 19» Some Representative Logarithmic Decrements of the Chromosphere Seen in the 1952 Eclipse by Athay, Billings, Evens, and Roberts ..... 150 20. Logarithmic Decrements for the Chromosphere Near a Height of 1000 km. Normal Regions . . . 153 21. Logarithmic Decrements for Excited Regions of the Chromosphere Near a Height of 1.000 km . . . 154 22. Observed Intensities of H3 as a Function of Height above the Solar Limb ...... 156 23. Observed Intensities of M4 as a Function of Height above the- Solar Limb ...... 157 24. Observed Intensities of H5 as a Function of Height above the Solar Limb ...... 158 25. Observed Intensities of H6 as a Function of Height above the Solar Limb ...... 159 26. Observed Intensities of H8 as a Function of Height above the Solar Limb ...... 160

2 7 . Observed Intensities of H15 as a Function of Height above the Solar L i m b ...... 161 28. Computed Intensities of H3 and H4 as Functions of Height above the Solar Limb ...... 173

2 9. Computed Intensities of I-15 and H6 as Functions of Height above the Solar Limb ...... 174 30. Average of Computed Intensities for the Balmer Lines as Functions of Height above the Solar Limb ...... o ...... 175 xlx LIST OF ILLUSTRATIONS — Pont inued FIGURE PAGE 31o The Average Balmer Decrement in the Chromosphere

at 500 llCni » 0 0 0 *6 66000600006 OOO I ^2 32 0 The Average Balmer Decrement in the Chromosphere a0 1000 lCm 60000600006000000 0. 1^3 3 3 . The Average Balmer Decrement in the Chromosphere

at 2000 1cm OOOOOOOO OOOOOOOOOO iSlj*

34o The Average Balmer Decrement in the Chromosphere at 4000 lem 00000 000060000 0.000 185 ABSTRACT

An extensive search of the literature was made for the intensity ratios of Balmer line emission spectra for astronomical objects. Studies of forty gaseous nebulae contained in twenty- six publications were found and examined. Average decrements were confuted for each nebula, and those were found tb be sufficiently homogeneous to permit our computation of an average all-nebulae decrement. Thirty-five studies published after 19^9 were exam­ ined for thirty-eight stars of eight different types. General inhomogeneities in the methods of obtaining and reducing of these data together with those in the resulting intensity figures precluded our calculating significant averages for these objects. Balmer intensities for the solar chromosphere appeared in nineteen publicationss and these were used in our computa­ tion of the decrement and slope of the decrement at three heights above the limb. These figures were then used to calculate an extrapolateds average decrement as a general function of limb height between heights of 500 and 7000 km. These calculated average decrements agree well with observa­ tions above 1000 km. xx xxi Other sources of extra-terrestrial Balmer emission are briefly discussed as are numerous suggestions for obtain­ ing high-quality data* Numerical values for typical decrements are presented in the Summary Section of each chapter. INTRODUCTION

Extensive hydrogen emission spectra have been measured by many observers during the last fifty years in astronomical objects as diverse as gaseous nebulae, certain stars, some galaxies, and particular regions of the sun. It will eventually be of interest to compare Balmer decrements from these objects with those obtained from high- precision optical laboratory spectra. This study, based upon an exhaustive search of the astrophysical literature, attempts to further that goal by fulfilling a three-fold purpose: 1) To amass relative Balmer line intensities observed in various emission spectra over the years and to give an extensive bibliography from the literature. 2) To provide some intrinsically useful criteria for judging the value of the data found, and to Indicate where better data are most needed. 3) To present, where possible, the most likely values for Balmer line intensities in the various objects and groups of objects examined, as based upon the literature to date. The following conclusions are reached in the course of this study: l) We find that although a large number of decremental studies exist in the literature and have been considered 1 here, the great majority of them are not highly accurate. Me hope that all available high accuracy data have been included here along with those from the large number of less accurate studies. Data have rarely been taken often enough or with adequate accuracy, 2) Me find several important requirements must be met in astronomical emission studies, if highly accurate and meaningful decrements are to emerge; a) the quantity meas­ ured must be proportional to the total photon emission in each line 5 b) the spectrum must be sufficiently dispersed and resolved in order to make clearly measurable the emission in as many lines as possible; e) corrections must be made for atmospheric effects (absorption, self-reversal, etc.) as well as for optics of the observing system and various photo­ graphic and photoelectric phenomena; d) the differential absorption effect of the interstellar medium must be included as a correction for distant objects near the plane of the galaxy; e) finally, the time-dependence must be carefully considered for objects with varying decrements, 3) Me find published data for gaseous nebulae and for the solar chromosphere to be reasonably eonsistant; and thus we have defined and computed here average Balmer decrements for a) the solar chromosphere at different heights; b) for each nebula; and e) for a typical "average" nebula. An attempt has been made to show to what extent various pub­ lished decrements vary from our averages. 4) We find published data for other astronomical objects, i.e.. Be stars, variable stars, galaxies, and solar flares and prominences, to be sufficiently incompatible with one another to make computation of an average or "typical” decrement meaningless. For these objects we present all available data and discuss the merits of individual studies. 5) We often find in the literature that details of observational techniques and data reduction techniques are either obscure or omitted. The reader of such articles is thus too often forced to judge results on the author’s general reputation or to look for agreement with other articles of equally unclear nature. 6) We find that photoelectric and photographic tech­ niques used in conjunction have yielded the best results in the field of nebular spectroscopy. Photoelectric techniques have recently been initiated for stellar observations. We offer a suggestion for the use of photoelectric technique for chromosphere measurements in eclipse = GENERAL REMAREE

The Definition

The Encyclopaedic Dictionary of Physics (Thewlis , 1961) says of the Balmer Decrement: As the study of the distribution of emitting atoms among the various quantum states5 i.e., for the calculation of the population of the excited levels, one examines for hydrogen the relative in­ tensities of the lines of the Balmer ... series. Observations provide the sequence of the ratios l(Ho): l(Hy): ... which is termed the Balmer decrement. This•sequence is then compared witiTTKe requirements of the cur­ rent solar and stellar theories... < This provides the definition of Balmer decrement for this study, and any tabulation or graphing of this intensity series will also be referred to as the Balmer Decrement. The line intensities themselves are geometrically normalized measures of source power proportional to the product of energy per photon and the photon flux in each Balmer line.

Problems of Measurement

The principal problem encountered in measuring these intensities is the accurate determination of photon fluxes. Ideally a photon detector should count every photon in the wavelength range of interest, which in our case is 1x3650-

6563!. In reality, however, photoelectric and photographic devices are the convenient ones to use* and they count only a small percentage of the photons striking them at any given wavelength. They also do not detect photons at all of these wavelengths with equal sensitivity. Several chapters in Hiltner (1961) discuss these detectors in detail3 a few of the major points are repeated here. Photographic plates should be calibrated individually to determine accurately how different wavelengths and inten­ sities affect image densities on the plate. The image results from an incompletely understood process by which individual photosensitive grains of silver halide are first acted upon by light and then chemically processed. Besides the low ratio of reactions to the number of incident photons* a number of non-linear factors make any direct measurement of photographic data difficult, "Reciprocity failure” is one of these. An exposure of x minutes to light of inten­ sity y does not give the same number of grains affected as exposure for y minutes to light of Intensity x. Several peculiar "adjacency affects” caused by diffusion of exposed and unexposed grains must also be accounted for. Photomultipliers must also be calibrated for their wavelength dependent response * but they are far more linear and accurate than photographic plates. They have a much higher ratio of measurable events to number of incident photons and a far greater dynamic range, Photomultipliers compare unfavorably with the photographic plate in some ways„ They can observe only a limited frequency band at one time which decreases the observational integration time; more­ over "toien they.are employed/in slitless.'operation their V relatively"large detecting surfaces/limit Instrumental resolution» Electronic noise often makes determinations of low-level intensities highly unreliable» We find in this study that the best features of these two observational techniques can be utilized coopera­ tively in single studies to yield the most accurate raw data obtained to date. Further comment on the relative merits of the two techniques is postponed until the main body of this thesis.

Corrections to the Raw Data

Even if one had an ideal detection device with which to make earth-bound measurements 5 he would still not be able to determine everything about the objects he was viewing. He is* in effect * looking at the output of a black box con­ taining a source and extending to his detector. In order to make physically meaningful statements * he must make succes­ sive corrections toward the source of radiation* effectively shrinking the black box and the unmeasurable effects which it contains« 7 Correctlen number one compensates for wavelength dependence of the optics of the system. This is a unique property of the instrument used. Correction number two compensates for atmospheric extinction5 that is, the wavelength-dependent absorption properties of our atmosphere. Intensities may be compensated for both this effect and the one caused by the instrumental optics by use of a standard star for calibration purpose. Such stars are considered to have known energy distributions. When one of them is chosen near the same zenith angle as the object being studied, its spectral distribution serves as the standard by which unknown intensities may be calibrated. Thus intensities of emission lines can be corrected to out­ side of the earth’s atmosphere. (To be truly accurate, the standard star should have its spectrum measured accurately from above the atmosphere by a system whose optical charac­ teristics are well known. We do not yet know of this having been done.) Correction number three is an adjustment of observed spectral intensities for the wavelength-dependent interstel­ lar absorption. In many instances, data calibrated to a point outside of the earth’s atmosphere are adequate; but this correction should always be considered. Shajn (193^) noted that the slope of the Balmer decrement decreased with Increasing galactic latitude. To compensate for the reddening caused by galactic dust, one searches for a general function f(X) such that

^"actual^) ” * observed^ +

The currently accepted f(X) curve resulted from studies of reddened and unreddened stars made by Whitford (1948, 1958) and by Wampler (1961). A schematic reddening-curve is pres­ ented in Fig. 1. Absorption (in magnitudes) is plotted as a function of inverse wavelength. This curve is valid for almost all of space; one of the exceptions, near the Orion nebula, is discussed by Mendez (1963). With f(X) specified the problem of obtaining lactual^ 13 reduced to obtaining the constant c for particular objects. There are two general ways by which one may attempt to do this: a) by assuming one knows the true energy distribution of some standard star in the region and comparing this spectral distribution with measurements; b) by assuming that the Balmer lines in a nebula result from pure recombination and then comparing the Paschen and Balmer line intensities from the same upper levels. If n is one such level and one considers degenerate states, then the ratio of the intensities of the Paschen-line to the Balmer- line from that state is 9

X („) ------

.35 .4 5 .7 1 2 =o

0

(magni­ tudes)

15

3 2 0 X

Figure 1. A Schematic Reddening Curve. 10

If Pn) W”3hvn3 HnA 3|lvn3 A 3vn3

using Herzberg1s (1944, p. 152) expression for the inten­ sities. W11111 is the number of transitions taking place per second in the light source from level n to level m, hvnm is the energy of a radiated quantum. Nn is the number of atoms in the initial state n, and Anm is the Einstein transition probability. The resulting ratio is independent of any physical parameters of the nebula. To be completely accu­ rate one must consider angular momentum states and selection rules. Since for a given level n not all of the Paschen and Balmer lines originate from the same 1-states, there is some dependence in the Paschen to Balmer ratio on the populations of these states. Kaler points out that this ratio is well known and not very sensitive to the electron temperature of the nebulae. The shape of f(X) and the use of the reddening curve are discussed at length by Sharpless (Strand 1963, chapter 12).

The Uses of the Balmer Decrement

After making the corrections discussed, an investiga­ tor has obtained Balmer line intensities as they would be measured emerging from the source; from these he can attempt to explain the physical state of the observed object. He gains information from the details of microphotometric trac­ ings of line profiles as well as from the intensities. He may attempt to give atomic level population distributions, atomic and electron densities, and various temperature measurements. He may try to account for the observed decre­ ments by suggesting that certain physical processes are occurring such as self-absorption of emitted light or non- radiative transitions. For examples, see Aller (1963) chapter 7 and Thomas and Athay (1961), pp. 280-352.

This study does not consider the further analysis of intensity data. Our purpose is to obtain the relative total radiation from each line. The success of published studies in presenting this data accurately has often been a function of the author's definition of intensity.

Types of Intensity Data Found in the Literature

l) The intensity integrated under the line profile 1 = J n m - icont(x)]dx shown in Fig. 2a is indeed the total radiation from the line. Intensity units disappear in taking ratios. Any quantities proportional or nearly proportional to these integrated intensities can serve as acceptable data to compute Balmer decrements. B W,’(x) w!(x)

X 2

Figure 2. The Three Common Measures of Emission Line Intensity.

A. Total Emission gives this area; I. = [l(X) - I (X)]d\ . i over the line B. The Equivalent Width of an emission line is the width of a rectangle whose height is determined by the continuum and whose area equals the area of the emission line above the continuum. C. The Central Intensity is the maximum intensity above the continuum. 13 2) Often the equivalent width of a line, defined as

TCKX) -2opnt(x)]dx w '( x t ) = ------cont shown In Fig. 2b, serves as a reliable measure of the line's intensity. Its units are "equivalent Angstroms," and it may be thought of as the width of a rectangular area under the level of the continuum intensity which contributes an amount of energy to the spectrum equal to that of the given line. Unless the continuum level underlying the Balmer line inten­ sities is far from constant, this is a good measure of relative intensities. 3) In the event that the Balmer lines are well in focus and have approximately equal line widths, it is reasonable to assume that the equivalent widths are nearly proportional to the central intensities. See Fig. 2c. In these cases it would be possible to compute relatively accurate decrements from the central intensities. If poor focus, self-absorption, or any other factor causes the lines to differ substantially in profile from one another this method is correspondingly inaccurate. All three measures of intensity occur in the litera­ ture, with the least accurate being the most common. 14 Normalizing the Data

In the subsequent discussion the common astronomical practice of setting l(H0) = I(H4) = 100 is followed„ The other Intensities are then normalized to this intensity to form the decrement. This practice is common because H(4) is the least often blended of the bright low-n Balmer lines and because it is relatively easy to measure its intensity accurately. This technique is also advantageous for com­ paring several decrements. GASEOUS NEBULAE

In this stndy» gaseous nebulae comprise the easiest group of objects to examine spectroscopically. Their spectra are by far the least complicated and are the slowest to change. They have been more frequently examined than any other distant group of objects of homogeneous nature and have been the most carefully and meaningfully examined. Perhaps the improving quality of observations for these objects is due in part to efforts of theoreticianswho have tried to derive the physical state of this relatively simple group of objects from spectral observations thereon. Following their attempts to formulate a correct description of nebulae and to account for previous anomalies, more careful observations have been made to point up new anomalies. This data-refinement process has continuously pro­ gressed until today a sizable literature of observations exists, but of highly nenuniform quality. This is discussed at length in the following remarks.

The first extensive spectral observation of gaseous nebulae were undertaken by Wright (1918), who simply made

15 16 eye"estimates of the intensity of the lines, Mierophoto~ metric tracings preceded awareness of photographic errors (Berman 1930)s and until the 1950's a set of intensities measured by one observer could rarely be correlated to another’s data. The first correction of the data for spec­ tral reddening due to differential interstellar absorption appeared in 1956 (Aller and Minkowski) and has to date been incorporated into less than a dozen publications. This is not to imply that this correction lacks Importance; it will* in fact* be shown that a correction for reddening which changes the intensity of the Balmer lines by at least 50$ is not unusual. The most recent* and apparently most reliable* results have incorporated complementary observations by photographic and photoelectric techniques* a treatment first used by Aller and Liller (1959)• The observational art has been refined to the point where Kaler (1964)* in comparing his results to theory* concludes that roughly half of the nebulae which he has examined have Balmer decrements which differ from theoretical predictions. Theory in its contem­ porary form* he stated* seems to predict the lower limit of the decrements. He found no correlation with the deviation of the decrements and any physical parameters* and has been able to rule out the following sources of deviation: 17 1» observational error 2„ overcorreetion: for Interstellar reddening 3. self“absorption In the Balmer lines 4. atomic eollisions 5« line emission from the central star (planetary nebulae) 6« expansion and thermal velocities of the nebulae 7, radiation from regions of different temperature. These results are cited here not because theory is being examined, but because attempts have been made to explain and, in some studies, t° correct, the observations by means of these parameters. The implication is that in order to obtain correct line intensities and the Balmer decrement, it is sufficient to measure nebular Balmer line intensities, and to make only a reddening correction consistent with the nebular galactical latitude and distance. These factors are also cited because they are very likely needed to account for the anomalous decrements observed in stars and galaxies which we consider in later chapters.

The Published Bata

In this section we tabulate all available data as well as giving notes concerning certain of the sources, an evaluation of the source articles, a weighting of the data, an attempt to sort the decrements into classes, and finally a computation of an average nebular decrement. For simple reference to the many articles involved, a system of convenient abbreviations is utilized for the most-used sources. The list of these follows: 18 Wr Wright (1918) Be Beraan (1930) P Page (1936) GH Greenstein and Henyey (1939) BW Bowen and Wyse (1939) A4l Alien (1941) Wy Wyse (1942) M-LD Mao-Lin and Dufay (1944) A51 Alien (19S1) WA Wllsen and Alien (1951) ABM Allens Bowen, and Minkowski (1955) MA Minkowski and Allen (1956) AM Allen and Minkowski (1956) Ge Geake (1956) Gu Gurzadyan (1956) Bu Btangess (1958) AL Allen and Liller (1959) S Seaton (i9 6 0) Ma Mathis (19 6 2) ABW Allen> Bowen5 and Wilson (1 9 63) LA Liller and Allen (1 9 6 3) 0 © ’Bell (1 9 6 3) M Mendez (1963) AKa Aller and Ealer (1964a) AKb Aller and Ealer (196413) K Ealer (1964) To evaluate nebular data, we first compare all the available data for that group of nebulae upon which one or more photoelectric studies have been made. These nebulae and the published studies thereof are given in Table 1. The underlining is explained later in the text. The following tabulation of published data is sepa­ rated for each nebula into the following divisionst 1 e photoelectric data, corrected for reddening— This may include articles in which all data are determined pheto- eleetrieally or in which selected photoeleetrioally measured lines serve to calibrate photographic data. The differential 19 Table 1 Nebulae Examined in Details with References Thereto

Nebula Studies Included in This Paper

NGC 1535 S LAS m , A51, A4l, ¥r m o 1976 & Wea Maa AL., Gu5 M-LDj, Wy5 QH3 Plaskett5 Mr m e 2 0 2 2 LA 5 B u 5 Geff MAS A4l m e 2 3 9 2 0„ LA5 S fl Bu 3 MA m e o 244 s LA, EAa A51, Wyfl Mr m o 6 5 4 3 0 , LAa A4la Bea Mr m e 6 5 7 2 K, AKa LAa A4la BWa Be m e 6 7 4 1 K. LAa A51s Mya Mr

N e e 6 8 2 6 o » LAa A4la Bea Mr m e 7 0 0 9 & AKa3 M a Mya A4la Bea Mr m e 7 0 2 7 & Oa LAa ABMa Sa Bua AMa ABMa Mya A4la BM,, Pj, Be m e 7 6 6 2 S Oa LAa Sa Bua MAa Mya A4la BMa Pa Mr i c 3 5 1 LA3 Bua MAa A51, Mr 10 4l8 Oa MAa Mya A4la Mr i c 2 1 4 9 & ^ a A51, Mr

I C 4997 & Oa LAa A4la P 20 absorption effect of the interstellar medium (reddening) is considered. 2, photographic data, corrected for reddening .3« Photoelectric data, not corrected for reddening 4. photographic data, not corrected for reddening fhe following symbols will be used in the tabulation of the data to designate the above categories: PE-.R - category 1 PG:R - category 2 PE - category 3 PG - category 4 Tables 2, 3, and 5 through 1? are arranged with the decrements given vertically. To the left of each table are the values of n, and the values of l(Hn) appear under the abbreviation of the publication in which they appear. Decre­ ments in category 1 are just to the right of the n-valuesj decrements in categories 2, 3, and 4 occupy successive posi­ tions to the right. Within each category the decrements appear chronologically with the most recent in the position furthest left. Comments concerning reddening calculations and the relative value of each set of data accompany the tabulation. M # 1535 Comparison of the two Kaler (1964) decrements shows that for this nebula the reddening correction is quite small relative to other errors. 21 Table 2 Balraer Decrements Published for NGC 1535

A4l Wr nations K K :M m A51

3 361 cO 4 100 100 100 100 100 100 100 5 51 49+ 60.0 51.6 49 55 92 6 22 21+ 36.4 28.8 21 31 48 7 55* 40.0 30 60 8 19 17.4 29* 21.8 21 9 6.5 5.95 7.7 10 4.8 4.28 11 3.8 3.47 12 3.15 2 .8 2 13 2.55 2 .3 0 14 1 .0 0 1 .5 0 15 1,40 1.25 16 1 .6 1.46 17 1 .0 0 .9

PE: R ” PE T g

■ taken from A51

(For a complete explanation of the tabular form used see text pp. 19-20.) m e 1976 The Orion Nebula is an excellent example of an ex­ tended nebula, Unlike planetary nebulae5 which comprise the majority of objects cited in this chapter, an extended nebula does not surround,a central star but occupies a much greater volume, Mendez’ (1963) data illustrate how the decrement varies from point to point in this object. It is of inter­ est to determine if the intensities published for each point functionally relate to the number of lines that are observed at each of the points. This is easily checked by consider­ ing the sets of averages shown in Table 3* No systematic relationship is noted and the general average was computed by weighting the intensity at each point equally. In Table 4 one can see excellent agreement among the photoelectric studies for NOG 197.6» as well as the irregular quality of the earlier photographic studies. The tabulation also indicates that the reddening correction is very important. m e 2022 Two sets of observations were made by MAS one at Lick Observatory and one at Mount Wilson. The figures given in Bu include a reddening correction for the average of the MA data. The two photoelectric studies indicate that the raw intensities from earlier photographic work were not accurate? consequently the Bu data suffer from photographic errors. Table So Averages of Mendez8 11 Point Observations of MSG 1976

# of points points point point points average of weight (no. lines 1 & 2 4,6,7,8,10 3 9 5 & 11 _ all points of pts taken n observed 10 ----g— 6 in average)

.3 264 289 265 295 320 288 11 4 100 100 100 100 100 100 11

5 49.2 47.0 4 5 .6 46.5 46.3 4 7 .1 11 6 28.5 2 7 .6 2 7 .6 2 8 .2 2 9 .6 2 8 .2 11

7 22.5 21.8 21.0 2 1 .9 8 8 15.1 I6 .4 1 7 .0 14.2 1 5 .9 9

9 9.4 8.8 9 .5 8 .9 8 10 6.7 5.7 6.0 7 It 5.0 5.0 2 12 4.3 4,3 2

ro w Table 4« Balmer Decrements Published for NGC 1976 - The Orion Nebula

M((the numbers correspond to points on the nebula) ' 2 . 3 4 5 6 7 8 9 10 11 la AL

? 289 257 270 265 300 320 288 268 300 295 262 288 4 100 100 100 100 . 100 100 100 100 100 100 100 100 100 100 5 . 45 48,7 49.8 45.6 48.5 46,0 4 5 .0 46.0 4 7 .0 46.5 48.7 46.7 ^3 %5 6 28 2 7 .4 2 9 .6 2 7 .6 2 5 .4 2 7 .6 28.0 2 8 .0 2 7 .0 2 8 .2 2 9 ,8 31.7 32 2 8 .2 24.0 2 1 .0 19.8 1 9 .7 2 3 .1 2 1 .9 21.0 24.5 2 1 .6 I 1 5 .5 14.7 1 7 .0 13.5 1 6 .0 17.4 17.5 14.2 1 7 .6 9 1 2 .3 9.3 9.6 9.5 8.0 9.2 8.8 9.0 9.4 13.2 10 9.1 6 .1 7.4 6 .5 5.1 6 .0 5.7 5.1 9.3 11 6.6 4.8 5.2 6.4 12 5.4 3.6 5.0 5.2 13 4,6 14 15 3.1 16 2.6 17 2.25 18 2.05 19 1.75 20 1.55 21 1.40 22 1.23 23 1.17 24 .93 25 .8 7 26 =93 27 .80 28 .69 29 .71 30 .62 31 .62 PEiE Table 4 (Cont„)„ Balmer Decrements Published for NGC 1976 - The Orion Nebula

M K 8 9 10 11 AL 350+ 346.6 375< 464.4 421.2 394.2 432.0 402.3 441.7 437-4 350 100 100 100 100 100 lop 100 100 100 100 100 100 100 4 1 1 40.8 41, 3 2 .6 40.1 41.1 3 6 .1 3 6 .1 3 6 .8 3 6 .8 36.9 34.4' 41 6 20.9 22 ,I 17.2 1 8 .9 23.9 1 9 .9 20.2 20.4 20.6 20.3 2 0 .7 25 1 7 .4 1 5 .0 14.0 1 3 .0 15-6 14.6 14.5 15-0 24.4 I 1 0 .9 1 0 ,2 9 >0 9-3 10.1 11.4 11.5 9-4 10.3 18.1 9 10.4 6 .3 6.6 5 ,0 5 -3 5-9 5-7 5-9 5-5 1 0 ,9 10 7.7 4 .1 5-1 4.4 3.1 3-4 3-6 2.9 7-8 II 5.6 3 .2 3-4 5-4 12 2 .4 3.2 4.4 B 3.5 ll 4, 2.54 II 2,14 17 1 >86 18 1 19 1 20 1 21 1 22 1 ,02 23 0.97 24 .89 25 .78 26 -73 27 28 : 8 29 .58 30 .59 31 .51 32 .51 PE ^Values from AL Table 4 (Cent *), Balmer Decrements Published for NGC 1976 - The Orion Nebula

Gu M-LD Wy GH Plashett Wr 3 500 35 Oi 500 400 4 100 100 100 100 100

5 41 60 35 60 48 100 6 19 40 20 37 19 60 7 13 28 20 20 40 8 15 20 9 7.2 12 15 10 4,8 10 10 11 4,0 6 6 12 3.2 6 3 13 1.6 6 2 14 7 15 ,4 6 1 16 ,8 6 2

ro o\ 27 Table 5» Balmer Decrements Published for NGC 2022

— Fublica- n tlon Bu LA Oe MA(Mt». W) MA(Llck) A4l T" - —J75— -TUcT — ------4 100 100 100 100 100 100 5 52 60 55*5 41.7 39.1 87* 6 27 3 0 .5 19.4 1 5 .8 44 7 30.1* 8.3 8.7 44 8 61.2* 98 PGsR PE PG *blend

m o 2392 The raw data given in 0 (see Table 6 ) are corrected for reddening. Except for H3 this correction is small. K corrected the data given in MA for reddening; this correct tlon was also small. The general observational results are in reasonably good agreement with each other.

Table 6. Balmer Decrements Published for M C 2392

n tion 0 K 0 LA S Bu MA 3. 275 302.^ 3 3 2 294 319 4 100 100 100 100 100 100 100 5 49.0 44.7 45.6 48.6 48.5 47 44.1 6 25.7 29.5 23.4 2 5 .0 32.5 34 28 7 41,7 36.3 36.9 8 24.0 I8.-6 21.4* 1 0 .1 17.4 .9 8.3 4.9 9.5 7.7 10 6 .0 6 .9 5.6 11 5.1 6.0 4.8 12 3.1 3.6 2,9 13 2.2 2.6 14 15 1.7 1.9 16 1.5 1.8 PE:R PGiR PE PG *blend m o 2440 The reddening correction f&r this high latitude object is '-.negligible *:: as may:- be seen in Table 7# "so all data must,be .'evaluated . The •.photoelectric data shall be considered mdst reliable, and, on that babis, we. believe the earliest studies, Wy and ¥rs deviate'-too greatly from later studies and tre do not include them in our average decrement for M O 2440.

Table 7 = . Balmer Decrements Published for NGC 2440 n tion IC K LA m A51 Wr 100 "Tocr^" 100 270 100 100 i o o : 100 100 100 100 5 52 ■ ::52-5- 4 3 .0 52+ 52 30 104 6 26 ' 26+ 26.3; 26+, 26 30 4l 7 45.9 15 49 8 20.4 20.4* 12.0 2 0 .4 20 9 9.7 •i:'9*7 r m 1 0 .1 10 10 7»7 7.7 4.4... 11 5.9 5.9 12 5.4 5.4 13 4.45 4.45 14 5.65 5.65 3.6 3.6 fi 3.7 3.7 17 2.95 2.95 18 2.35. 2.35 19 2.35 2.35 20 2.35 2.35 21 2.15 2.15 22 1.8 1.8 23 1 6 j 1.3 PE:R PE PG 4- from A51 ^ from MA 29 NGC 6543 The reddening correction as shown in Table 8 alters uncorrected Intensities by about 20$. The earliest data# Be and Mr# deviate greatly from the photoelectric data.

Table 8. Balmer Decrements Published for NGC 6543 Publiea- n tion 0 0 LA A4l Be Mr 3 “316 4 100 100 100 100 100 100 5 5 0 .1 45.7 60.4 41.3 3 6 .3 96 6 3 0 .2 2 5 .7 30,7 21.4 1 3 .8 59 7 33.1 2 7 .5 31.7 5.4 59 8 1 9 .5 1 6 .2 21.4 53 2.2 . 35 9 7 .8 6,3 9.2 1.1 32 10 3 .0 2.4 6.4 0.5 16 11 6.9 11 12 6,7 PB;R PE PG

NGC 6572 The reddening correction for NGC 6572 is quite importantt it increases the observed Intensities by 50 ~ 100$. The reason for the steepness of the LA decrement with respect to AKb and K is not understood. The,photo­ graphic intensities in A4l are comparable with later data# but decremental agreement of Be with the other measurements is weak.

Reddening distorts the observations of NGC 6741. Other than the blended lines observed# n = 5 and n = 8# the 30 Table 9. Balmer Decrements Published for HOC 6572

Publiea~ n tion K K AKb M A4l BW Be

3 251 370* 500 4 100 100 100 100 100 100 100 5 65 31.0 34.8 5 6 .6 3 6 .0 55 6 27 21.9 2 7 -6 26.8 2 1 ,1 31 7 39,8 42,7 35 23 8 24 1 8 .1 1 8 .1 14.3 14.3 16 9 9,8 7.3 7.3 8 .5 6 ,0 13 10 7,4 5,4? 5,47 5,1 3.4 8 U , 7,3 5.4 5.4 3,9 2 .3 7 12 6,2 4.5 4.5, 2.8 7 13 3.6 2.64 2.64 14 3 .8 . 2.72 2.72 15 2„9 2.10 s.io 16 2.45 1.70 1 .7 0 17 2.15 1.51 1 .5 1 18 1,60 1.14 1.14 19 1,38 0 .9 8 0 .9 8 20 1.25 0 .8 9 0 .8 9 21 1,10 .77 ,77 22 1,02 .73 »73 23 087 o62 ,62 ,87 ,62 25 »71 »50 26 ,78 ,55 ,55 27 ,62 ,44 ,44 2 8 ,5 6 .40 .40 29 , 6 0 .,.. ,43 ,43 5i " «36 .36 31 *54 : ,3f ,.38,, 32 .51 .36 ,36 33 .54 .38 .38 34 35 _ _ _ _ _

PEjR • PE PO ^Aller unpublished photoelectric data sets are quite consistent. The early . studies5 Wy and Wrs are here: in reasonably close agreement with later work.

Table 10,„ Balmer: Intensities Published for NGC 6741

— mbiica=— — ------— ------— — —

n tion K K •LA . 451 Wy Wr i> 470 4 100 100 100 100 100 100 5 48 374 - 5 2 .6* 37 35 59 6 42 3 CH- 30.5 30 25 32 7 24.8 12 42 35 8 1 7 .2 9.0 7-4 12 9 11 7-1 10 9-3 5-75 11 7 -8 4.82 12 4.3 2 .5 6 13 3-6 2.20 14 8.3 5 -0 0 15 2 .6 1 .5 6 16 1.65 1.00 17 1 .1 0 0 .6 5 18 0.7 0.4

PE;B ' PE PG v-f from A51 * blend

NGO 6826 Correction for interstellar reddening is slight for

NGC 6 8 2 6. The photoelectric studies agree well for the early Balmer liness but the photographic work is not in good agreement with the photoelectric. 32 Table li» Balmer Decrements Published for NGC 6826 E%bllca- "...... : 7' ‘ r ■ n tion @ 0 LA" A4l ' Be Wr 2oo T 100 100 ' 100 100 100 44.? 47.3 40 24.5 104 1 24.0 25.1 19 14.8 61 30.9 3 0 .6 6.3 63 I 17.4 23.5^ 51* 3.2 9 6.8 14.5 2.0 10 4.5 7.6 PE: R PE PG * blend m e 7009 The reddening affect was considered negligible by Kaler (see Table 12). The photoelectric sets of data have good agreement. The earlier photographic stmdies are of varied quality. m e 7027 This group of data is especially interesting. The reddening corrections made by K and 0 are made upon their own datag but the decrements given in three articles3 S 5 Bu3 and.AM* all result from correction attempts made upon ABM. The AM attempt was the earliest of these and seems to be the crudest. The agreement between the most recent studies 9 K* Os LA, ABWs is quite good. These also agree well with the 1955 study* ABMs upon whose carefully measured decrement the earlier differential absorption corrections were made. The earlier measurements are of questionable worth. 33 Table 12 „ Balmer Decrements Published for NGG 7009

Publican n tlon K K LA AKa My A4l Be Mr 3 288 288 470 145 4 100 100 100 100 100 100 100 5 49.1 49.1? 49.1 25 34.2 56 100 6 26.5 26.5? 28.3 24.8 20 24.9 35 60 48,4 46.5 10 39.5* 22 84 I 1 8 .9 18.9? 18.4 19-3 12 18.5* 14 12 9 6.15 6.15 8.8 6.10 8 9 10 10 4.28 4.28 5.2 3.76 7 16.6 11 3.38 3.38 3.8 2 .7 8 12 2.43 2.43 3.4 3.83 7 13 1.94 1.94 14 2.32 2.32 15 1.27 1.27 16 1.17 1.17 17 0.99 0.99 18 .82 .82 19 .68 .68 20 .60 .60 21 .49 .49 22 .46 .46 23 .42 .42 24 .3 8 .3 8 25 .31 .31 26 .32 .32 27 .27 .27 28 .2 7 .2 7 29 30 .24 .24 31 .21 .21 32 .21 .21 33 .22 .22 34 35 PE PE;R PE P§ f originally published In AKa 1964a. * blend Table 13. Balmer Decrements Published for NSC 7027

T u b H c a tion AM K AB¥ ABM Wy A41 BWP Be 65o~^T 5 o ”5Ho~ 100 100 100 100 100 100 100 100 100 100 100 100 100 100 100 49.0 46.8 48 49 40 33 30.9 32.9 33 20 2 7 .1 39 46 36 I 25.1 26.3 2 9 .0 29 21.3 14.0 14.1 15-4 I? 16,3 12 1 3 .0 17 10 19 7 52.5f f 2 5 .0 24.4 26.4 24.4 8 2 1 .6 14 15 12 16.6 20.4 18.5 12.6 8.0 9.1 7.9 2i 8.8 7 8 .0 13 8.6 9 11.2 9.33 10.1 6.6 .2 3.9 4.7 5.2 4.6 3 3.1 8 3.3 6 1.0 10.2 7.9 5.2 I.5 2.9 4.5 3.5 2 6 5 11 7.1 7.6 3..2 3.2 3.3 1.5 1 .2 4 3 12 5.3 6.3 2 ,.8 2.8 2.7 1 .2 5 13 6.2 2.6 2 .6 2,3 14 3.7 3 .7 2 6 15 5.3 8 2.2 2 ,2 5:1 16 4.4 3.4 1.8 1 .8 1. 17 4.2 1.7 1 1. 18 3.9 1.6 d 1.3 19 3.7 1.5 1.5 1.l to 2.7 1.1 1.1 1 .0 21 2.5 1.02 1.02 0 9 22 2.45 0 .9 8 0.98 0 9 23 2.35 0.93 0.93 0 .8 24 2.45 0.98 0.98 0.6 25 2.8 1.13 1.13 26 2,6 1.02 1.02 27 2.3 0 .9 2 0 .9 2 28 1.95 .77 0.77 29 1 .7 0 .67 0.67 30 2.04 0.82 0.82 PGsR "W W t blend 0 * taken 1 Table 14. Balmer Decrements Published for HOG 7662

”m H e at," ■ LAMA n tien :k 0 S . % E 0 Wy A4l m P Mr 157 229 "W 278 3301^““" W ™ 39 0 3o0 ■W"l; 500 h loo 100 100 100 100 100 100 100 100 100 100 100 100 5 46.8 47.9 51 50 43f 40.7 43.6 44 4o 41.7 40 3 6 .0 108 6 22.4 28.2 2 8 .7 28 19.3 22.4 23.5 2 3 .2 25 1 3 .0 25 17 51 7 17.4 41.7* 14.7 3 1 .6 34.9* 17 20 21.6 80 8 .3 71 8 17.8 17.4 14.7 12.6 7-8 7-7 25 8.0 25 52 9 8.72 6.03 1 0 .3 7-02 4,37 7.1 5.7 10 3-1 8 . 10 10 7.25 3»39 7-8 5 .8 6 2.40 4.6 8 11.1 6 8 11 5.90 6,3 4.79 7-7 3.8 5 1.2 5 12 4.90 5-2 4,02 3-4 4 1.2 4 13 4.46 4.7 3 .5 8 3-6 4 4 14 4.57 3-66 2.2 5 5 15 3.10 3-1 2.46 2.3* 3 0.5 3 16 2.76 2.19 1.5 17 2.45 1.97 1.3 18 2.19 I .72 0.9 19 2.00 1 .6 0 0.73 20 l.jo ,35 21 1.41 22 1.38 1.1 23 1.10 0.06 1.00 .79 0.85 .67 ..91 ' .73 27 .63 28 :p 061 29 .6? 30 53 31 .62 — w m — “* P©: R P© f Aller ) mpubllshed * "b lend w Bedtienimi; is seen to play a relatively minor role for HGC 7 6 6 2= The raw photoelee'ferie Intensities agree well with one another and with the more carefully done of the photographic studies5 particularly with MA from whose obser­ vations the corrected decrements of S and Bu were calculated, K and 0 again made eorrectiens upon their own observations« 10 351 The data given in Bu (see Table 15) are a correction of the decrement in MA» E did not find it necessary to correct at all for reddening. All of the data included here are worth considering further with the exception of the very early Mr eye-estimates.

Table 15. Balmer Decrements Published for 1C 351 "PuB33JeS n tion K Bu K M MAA A51 Mr 3 4 100 100 100 100 100 100 100 5 50 52 50t 41.6 49 58 89 6 30 27 30t 24 23 49 42.3 17 66 I 86.0 6.1 28 9 9.90 9.90 6.2 19 10 6.91 6 .9 1 9 11 4.S3 4.83 123»©3 3.83 13 2.48 2.48 14 I .7 4 1.74 15 1.74 1.74 .2.6 1,88 PBsR PQ:R HE PS t from A51 37 Table 16. Balmer Decrements Published for IC 418

Publlca- n tion K 0 K 0 Wy A4l Mr 3 282 282 398* 398 410 4 100 100 100 100 100 100 100 5 45.7 46.8 40.7* 40.7 25 45 96 6 24.6 2 5 ,2 20.2* 20.4 20 24 57 7 15.5 18.6 12.6* 14,5 15 39 8 . 14.5 15,1 11.6* 11.5 15 12.6 35 9 6.8 7.41 5.30 5.50 6 6,4 33 10 5.5 4.47 4,25 3.31 4 3.4 24 11 5.0 3.85 4 18 12 4.2 3.16 16 13 3.6 2 .7 0 14 3.6 2.72 15 2.75 2.09 16 2.35 1.77 17 2.15 1.61 18 2.15 1.63 19 1 ,6 0 1.22 20 1.50 1.12 21 1.40 1 .0 6 22 1,10 0.82 23 1.05 .78 24 0 .9 6 .71 25 . . 98 .72 26 .85 .63 27 .93 .69 28 .71 .52 29 .68 .50 30 .6 0 .45 31 .40 .36 32 ,51 .38 33 .42 O .31 PE: R ' PE PQ * from 0 38 1C 418 The reddening correction amotmts to 25 or 3C^ in 1C ;4l8. The photographic studies of My and A4l (gee Table 16) correlate fairly well with the photoelectric work. The data from Mr is incommensurable with the later work» ,10 2149 K and 0 both corrected O ’s 1963 data for differential interstellar absorption and found a much shallower decrement. Interestingly, the A51 photographic uneorrected data almost agrees with the carefully corrected photoelectric data.

Table 17° Balmer Decrements Published for IC 2149

Pubiica- n tion K 0 0 : A5i Mr 3 219 219 269 4 100 100 100 100 100 5 49.1 49.0 39° 8 52 100 6 ■ 25 »2 25.7 18.6 30 74 7 26.3 17,8 15 58 8 22.4 14.8 12 44 9 9.1 12.0 6.0 9.2 20 10 5 = 4 5°7 3,5 9.0 16 11 5°2 7 .3 12 12 4.3 __ 8 . PE;R ' PE PG x g m i Here again K and 0 find that appreciable reddening has occurred for a nebula they examined» The pre-1960 decrements, A4l and P, are of questionable accuracy in this case. Table 18» Balmer Decrements Published for 10 4997

Publican n tlon it 0 K 0 LA A4l P 3 246 257 427 4 100 100 100 100 100 100 100 5 55.0 51.3 45.0 42.7 51.0 35 56.8 6 30.2 30.9 22.4 22.9 25.5 17.8 14.3 7 46.8 33.0 32.4 3 8 .8 22.8 8 1 7 .0 14.6 11.5 16.7 15.3 14.2 9 7.25 6.0 4.63 4 .0 9.0 9.2 10 5-9 5.8 3.78 3 .6 2.7 5.3 11 5.3 3.38 12 4.3 2.69 3.8 13 3.3 2.06 14 4.6 2,81 15 2.6 1.62 16 2.1 1.30 17 1.75 1.08 18 1.35 0.83 19 1.1 ..70 20 1.1 .67 21 0.85 .53 22 .79 .49 23 .72 .45 24 .68 .42 25 .62 .38 26 .56 .34 27 .51 .32 28 .50 .31 29 .47 .29 30 .46 .28 31 .42 .26 32 .41 .25 33 .44 .27 PEiR .. PE PG 40 of the Published Data

A simple comparison of data of the type made on the preceding pages brings out several points; 1 o Although the reddening correction for a particular nebula.may be small or even negligible5 without considering it one cannot estimate reliably the worth of intensity data. One may determine if a reddening correction is necessary by noting the angle of the nebula above the galactic plane,. If the object is high5 a correction should be minimal. 2. fhe best agreement among sets of data has occurred for investigations utilizing photoelectric calibration. 3. In generalj, studies made prior to 1950 are unreli­ able. Later decrements are generally more consistent with one another. This is reasonable in view of continual refinements in measuring technique. For purposes of averag­ ing data to obtain representative Balraer decrements for nebulaes we have neglected these earliest studies,

. . . \ ■ ' Average Decrements for Each Nebula

Of the eighteen articles published after our 1950 cut-off date5 six have not been used in computing average decrements. These sl% did not include an interstellar differential absorption correct ions, and it was not made clear In any other article that such correction for the nebulae involved was not needed. The twelve remaining . 41 studies are - listed* ■; ■' K - Eller’s (1964) dissertation iss aoeording to Aller (unpublished remark made to author)s the most definitive work existing at this,time*. The decre­ ments of thirty-four nebulae are given5 fifteen of which are high-accuracy measurements made by com­ bining photoelectric and photographic techniques. In the photographic portion of his work many of the earliest lines are burned out on his plates. , Observational data are presented and corrected for reddening effect.

0 - O ’Dell’s (1963) work is of similar qualitys but is less extensive in that the entire study is photo- elecfrieally made with a scanning spectrometer. The intrinsic problems involved with this type of instrumenta i.e., photoeleetrie tube noise and resolutiona are his limiting factors. Both reddened and unreddened data are given. Me -Mendez (I963) has examined the Orion Nebula with 'considerable care, using photographic and photo­ electric techniques to best advantage. Evidence for considerable dust was found in the nebula. The reddening correction for this region does not take the St ebb in s-Whitford curve forma and the correction was made by comparing the ratio of Paschen to Balmer line strength ratios with theory. From his published point"by-point intensities5 it has been attempted here to calculate a reasonable decrement integrated over the entire object*

Ma - Mathis (1962) made a very limited number of photo­ electric measurements and attempted to see if the observed intensities were better accounted for by reddening or self-absorption and recombinations. The former fared better and his few observations and their reddening-corrected intensities are - presented here. AL - Aller and Liller published this early photoelectric

study in 1959 for the same extended object studied later in detail by Mendez * The reddening correction was also based on the ratio of Pasehen to Balmer intensities originating from the same hydrogen level. ' . S ~ Seaton (i960) presents data from other sources» MA and ABMj, and his corrected intensities after con­ sidering reddening. These data should have the same order of accuracy as the original measurements had Bu - Burgess (1953) had previously attempted a similar, correction upon the decrements in AM and ABM. Me also considered his corrected data to be as accurate as the observational data. AM - Although Aller and Minkowski (1956) made careful photographic observations, later studies indicate that their attempted reddening correction was not highly accurate. It was the first such attempt found to be published, and comparison indicated that subsequent studies have superseded it.

Only the previous eight studies considered reddening corrections. Because some nebulae discussed in these publi­ cations are cited as not having appreciable reddening, it is appropriate to include in our calculation of their average decrements data from articles not considering differential absorption effects. There are four such additional studies: AKa - Aller and Kaler (1964a) published along with their own data (which is considered under K with redden­ ing correction) the unpublished study of Aller and Faulkner for the bright ring nebula NGC 7009• LA - Liller and Aller (1963) had only limited observation time (which may have affected their accuracy) for their photoelectric measurements of line intensities in the twenty-five nebulae they examined. Other­ wise, they were limited only by the inherent photo­ tube difficulties discussed earlier. 44 MA “ This 1956 study by Aller and Minkowski was done photographically5 but was made carefully^ as the excellent agreement with the-later photoelectric efforts indicates» A51 - The remarks for the above study apply also to this 1951 paper done by Aller,

To compute an average decrement for each nebula, arbitrary weighting factors are selected for each publica­ tion on the basis of comparative accuracy of its measure­ ments and corrections. Indication of how relative quality is judged has been presented by way of the remarks immediate­ ly preceding and accompanying the data.

Table 19• Weights Assigned to Publications in Order to Compute the Average Balmer Decrement for Each Nebula

Ar¥icXe WelgEF assigned '" ' ' " ' " ...... K 3 (n < ? ) , 4 (n > 7)s 2 (data from other publica­ tions which K corrected for reddening) 0 4 (n < 7), 3 (n > 7) Me 5 for all n Ma 2 for all n AL 2 for all n S 1 for all n Bu 1 for all n AM 0 for all n AKa 2 (when correction for reddening is negligible)s 0 (otherwise) LA 2 (when correction for reddening is negligible), 0 (otherwise) MA 1 (when correction for reddening is negligible), 0 (otherwise) 1 (when correction for reddening is negligible) 0 (otherwise) 45 The articles used in computing the average decrement for each nebula are underlined in Table 15 page 19, Aver­ ages and total weights of the averages (the sum of the individual publication weights) are given in Table 20* Table 20 contains decrements for the nebulae which have Just been examined in detail. Table 21 contains aver­ age decrements for other nebulae listed in the eight publi­ cations already mentioned which made corrections for reddening. The twenty-four nebulae in this second group appeared in too few publications to be dealt with in the detailed manner afforded the first group* The raw data for the second group and the sources of these data are given in the first appendix.

Steepa Average s and Shallow Decrements

With average Balmer decrements tabulated for so many nebulaea the question arises as to whether these decrements may be grouped with respect to some intrinsic parametera for examplea their steepness. What.we mean by steepness is easily seen in Figure 3= For simplicity we have represented the decrements as straight lines5 although the actual curve is more complex« It is possible for us to find the average l(Hn ) for the nebulae and to determine if intensities for particular nebulae do diverge from the average as shown9 while n increases. 46 Table 20; Weighted Averages of Balmer Decrements for Selected Nebulae (First GroupD

NGC NGC NGC ; NGC' : NGC ■. NGC ■ NGC NGC NGC . NGC NGC NGC 1C 1C .10 10 1535 1976 2022 . 2392 ' 2440 6543 6572 ' 6741 6826 7009 7027 7662 351 418 2149 4997 3 2.84 175 286 245 .251.' 229 228 ■ 272 . 248 294 ; 282 . 219 252 4 100 100 •100 . 100 100 100 100 100 100 100 100 100 100 ; 100 100 100 5 51 46.3 .52 43.3 46.0 50.1 65 48 49.0 . 49.1 47.9 48.1 50 46.4 49.0 53.3 ■ tj 6 % 22 28.8 27 27 .4 26.2 30.2 ■ 27 42 27.5 26.4 2 7 .6 26.3 28 2 5 .0 25.5 30.6 8 19 I5.9 • 20.9 14.8 19.5 24 20.9 18.9 18.3 17.6 14.7 22.4 17.0 9 . 6.5 11.3 8 .3 9.0 7.8 9 .8 11 8.3 6.8 10.4 7.9 9.90 7.0 10.9 6.7 10 4.8 7.7 6.0 6.6 3.0 .7 .4 9.3 5.5 4.25 9 .7 5.88 6.91 • 5.8 •5.5 5.9 11 3.8 5.8 5.1 5.9 7 .3 7.8 3.25 7.2 5.98 4 •.'■83 5.0 5.2 5.3 12 3.15 4.9 ’ ■ 3.1 5.4 6.2- "• 4.3 3.20 5.5 4,96 3.83 4.2 4 .3 4.3 13 . 2.55 • 4.6 2.2 .4.45 3.6 3.6 1.94 6.1 4.52 2-.4i3 3.6 • 3.3 14 1.00 3.-65 . 3.8 8.3 2.32 4.3 4.57 1.74 3-6 4.6 15 i.4o 3.1 .1.7 3.6 2.9 2.6 I .27 5.1 3 .1 0 1.74 ' 2.75 2 .6 16 1 .6 2.6 .vl.5. 3.7 2.45 1.65 1.17 4.2 2.76 1.88 2.35 2.1 17 1.0 2.25 2.95 2.15 1.10 0.99 4.2 2.45 2.15 1.75 18 2.05 ' 2.35 1.60 0.7 .82 3.9 2.16 2.15 1.35 19 1.75 2.35 1.38 .68 3.7 2.00 1.60 1.1 20 1.55 2, 35 1.25 .60 2.7 1.70 1.50 • 1.1 2, 1.41 1.40 21 1.4o 15 1.10 .49 2 .5 i O .85 22 1.23 1, 8 1.02 .46 2.45 1.38 1.10 .79 1.07 1, 3 0.87 .42 2.35 1.10 1.05 .72 12 0.93 .87 .38 . 2.45 1.00 0.96 .68 25 .87 .71 .31 2.8 0.85 i • .98 . 62 26 .78 ' .32 2 „ 6 ,91 .85 • 56 27 :lo .62 .27 2 .3 .79 .93 .51 28 .69 .56 ,27 1.95 .78 t .71 ,50 29 .60 1.70 .85 .68 .47 30 -M .24 2.04 .66 .60 .46 31 .62 .21 .62 .40 .42 32 .21 i' .51 ..41 33 ill .22 t . .42 .44 weight for n>10 11 . 3 ■' 4 7 9 9 4 ■ 7 7 7 weight V 3 4 (5 to 4 (5 to .4 4 for n<10 n = 16) n = 15)

*Line 7 has been omitted because it usually blends 47 Table 21. Weighted Average of Balmer Decrements for Selected Nebulae (Second Group)

Anon Anon Anon Anon Anon NGC NGC NGC NGCNGC NGC NGC NGCNGC NGC NGC NGC IC IC IC IC J j BD l8h15m l8h47m 21h31rT1 22h29m 23h24m 40 2371-2 3242 4361 6309 6720 6790 6803 6807 6833 6884 6886 2165 4846 5117 5217 320 900 +30*3639 3 339 209 302 227 246 4 100 100 100 100 100 100 100 100 100 100 100 100 100 100 100 100 100 100 100 100 100 100 100 5 51 46.8 49 52 49 48 49.0 63 48 44 44 50 48 49 44 46 51 46 49 45.7 54.5 48 48 46.8 6 26 25.7 24.6 35 26 28 33.9 32 26 29 28 23 25 27 29 26 25 26.5 27 26.3 29.6 30 25 27.6 *8 11.5 25 19 18 14 14 18 18 21 14.5 16.5 20.9 9 8.1 8.1 6.2 10 11 13.5 7.4 11 5.8 14 5.9 15 7.2 8.9 17 10 6.3 8.8 4.1 6.0 9.6 7.94 10 5.0 7.1 4.8 9.0 8.9 10 6.5 8.7 5.3 9.5 4.7 10 5.6 6.9 8.5 3 .4 4.45 7.5 4.7 5.2 6.6 6.17 11 4.7 4.3 3.5 6.2 7.4 7.6 5.4 5.8 4.3 5.6 3.6 5.3 4.2 4.4 5.6 2 .9 3.7 4.1 4.3 3.2 5.5 12 4.7 3.9 2.95 5.4 7.4 6.3 4.5 3.9 2.6 4.5 2.3 3.7 3.8 2.9 4.9 3.3 2.9 4.9 2.7 2 .9 7.3 13 5.5 9.1 3.5 3.0 2.4 2.45 14 2.6 2.6 2.15 3.9 3.2 1.3 3.0 3.2 15 2.2 2.0 2.45 2.3 1.1 1.8 16 1.52 1.85 17 1.38 1.6 1.4 18 1.03 1.48 1.35 19 0.91 1.32 1.1 20 0.91 1.2 1.0 21 0.83 0.91 0.80 22 0.83 0.81 0.80 23 0.78 0.70 24 0.81 - 2 2 3-4 2 2 2 3 3-4 3-4 2 2 2 2 2 2 2 2 3-4 2 2 4-2 2 2 4 *Line 7 has been omitted because it usually blends. 4 n

Figure 3. Schematic Representation of Shallow, Average and Steep Balmer Decrements.

A. shallow B. average C. steep 49 We have averaged the intensities for nebulae group nebulae group 2«, and all nebulae and have tabulated these averages along with the total relative weight for each averaged intensity in Table 22. Figure 4, p. 51 j, shows a plot of the all-nebulae average intensities plotted vs. n, as well as plots for the individual nebulae. To simplify determination of the behavior of the decrement for each nebula# we defined an average group of intensities for each n# which included approximately half of the nebulae exhibiting that Balmer line. This group is indicated by the hatchings in Figure 3- Table 23 indicates the actual intensity ranges and the numbers of nebulae whose intensities were above# within# or below these ranges for each h. We have crudely determined the character of the decrement for each nebula by tabulating how each value of l(n) compared to the average intensity groups. The results may be seen in Table 24# where 0 indicates inclusion in the average group# and + or - exclusion# above or below respec­ tively, If a nebula has a column of zeros under it# its decrement may be said to be, average. If a 4- is entered for n = 3 and a column of - 1s extends down from n = 5 to the last visible line# the decrement is steep. Conversely# if - is entered for n = 3* and +*s fill the column downward from n = 5s the decrement is shallow. This characterization is 50 Table 22» ’’Average” Balmer Decrements for Gaseous Nebulae

Group 1 Group 2 All Nebulae n Weight I(Hn) Weight l(Hn) Weight I(Hn) 3 79 264 13 257 92 263 4 100 100 100 5 90 47*2 57 47*6 147 47*4 6 90 27*7 57 27*7 147 27*7 I 77 18*3 30 17*6 107 18.1 9 89 8.9 54 8.5 143 8.8 10 82 6.29 53 6.68 135 6.44 11 53 5*66 54 4.77 107 5*21 12 53 4.46 50 4.08 103 4.28 13 49 3*68 18 3*9 67 3*74 14 40 4.09 24 2.65 64 3*55 15 49 2.74 18 1.89 67 2.51 16 48 2.54 6 1.74 54 2.45 17 39 2.13 10 1.48 49 2.00 18 35 1*93 10 1*34 45 1.80 19 31 1.86 12 1.11 43 1.65 20 33 1.56 12 1.04 45 1.42 21 31 1.44 12 O .85 43 1.28 22 31 1.31 12 0.81 43 1.17 23 31 1.50 8 0.74 39 1.34 24 27 1.06 4 0.81 31 1.03 25 27 1.01 27 1.01 26 27 1.02 27 1.02 27 27 0.92 27 0.92 28 27 0.80 27 0.80 29 24 0,84 24 0.84 30 27 0.73 27 0*73 31 23 0.44 23 0.44 32 15 0.41 15 0.41 33 15 0.40 15 0.40 51

300

100

510 15 20 25 30 n Figure 4. Nebular Balmer Line Intensities as a Function of n . The circles represent the average intensities of all nebulae studied. The shaded area contains £ of these nebular intensities. Those intensities not included are distributed almost equally above and below the shaded section, cf. Table 23. Table 23» The "Average"- Intensity Range In Gaseous Nebulae For Each Balraer Line and the Number of Nebulae Whose Intensities Fell Within and to Either Side Of this Range

n Range of Intensities Lower Than Within Higher Than 240 - 4 8 4 100 0 35 0 45.80 - 49,0© 5 19 12 1 26.20 - 29 .2 0 12 18 6 15.30 20 .9 0 6 13 4 9 7.00 10.60 a - . 17 9 10 4.80 а . 08 6 19 9 11 4.3 0 б,12 8 18 6 12 5 .06 10 16 - 6 13 1 v i e 4.6 0 6 9 3 14 2 .6 0 4 .5 0 5 10 3 1, 91 3 .10 6 10 2 i i 1, 85 '3.05 5 2 17 1, 38 2.62 3 • 2 18 1, 35 2.25 3 7 2 19 1.1 2.2/ 2:' 7 2 20 1.0 1,84 2 7 2 21 0,85 1.71 3 6 2 22 0.81 3 6 2 23 O .87 4 5 1 24 0.87 1.19 3 4 1 25 0.71 1.31 2. 4 1 26 0,78 1.26 2 4 1 27 0.62 ‘ 1.22 2 4 1 2B 0.56 1.04 2 4 1 29 0.68 1.00 2 1 30 0.51 0.95 2 1 31 0,34 0.54 1 3 2 32 0.31 0.51 1 3 0 33 0.35 0.45 1 2 1 53 Table 24. A Comparison of Nebula Balmer-Line Intensifies With the Group of "Average" Intensities

Group 1 Group 2

6 s E S $3 cn LT\ b- H C^^t CM on H CO CXi CU 8 VO lA VO OJ OJ O CO CVJ H xo ^ CJ OV l> X5 jC X5 S h vs. oj H w, ch ^ ^ o o on g\ ^ cn^rin ^vo is - s_____ , on CO E~- cy ov F-CVJ O CU VO HCO^j-OveOCOHOJCO is- F ^-V VOvo Oo cmcvj cr\OvOOCV)OOCOVO^tHHOOo o o m QO VO r~l sH O O o OjlAOvO CO-d- LA iAt>COOOVO Lf\H HCAHHCVIOUCVlO on(OOJCOCOh-^-OQOOCOCOCOHCOHCy. ^ ^ (MOO H H HOJCVICXJVOVOVOVO^t^-h- CO^t OJ -=j" ^-OJCO^tVOVOVOVOVOVOVOVOCVl^lAtACOOVCOCM ^ ti G ^ S fl •4- ^QOOOOOOOOOOOO DO GO OQOOOOO OOOQOO o O O O P cb O OOcbcbcbOoo o o d ti do OOC5C3C3 CDC5CDOO OOOO O h ) h > A n % ^ a % ^ H H < <<12; K iz; S S id i a S H H H H p cn p 3 0 - 0 00 -+00+G-0 + — 4’ — 0 4 000000000000000000000 0 00 0 0000000 0: 000000 3 4-0 -!-=’0-!--t-00+00+00-f-t00+000+0---!- 00-04 00-+-000 6 — 000040400000 — 04 — — — 4 — 044 — 00 — — 00 — — 00 0 4 4 — 0

o o 0 — 0 4 0 0 0 0 0 4 0 0 0 Q: 4 4 9 — 4 0 0 0 0 4 0-000 0 4 0 0 0 4 4 04 - 0 0 0 0 —■ 0 0 0 10 0 0 0 0 - 0 4 0 — 4 0 00 0 0 0 0 0 4 4 4 04 0 0 0 0 0 0 “ — 0 0 0 11 - 0 0 4 4 4 0 0 0 0 0 0 0 -4 4 4 0 0 0 0 0 0 — 0 0 — 0 0 12 — 0 — 4 4 0 — 4 0 0 0 0 0 0 0 -4 4 4 0 0 — 0 — — 0 — 0 — 0 4 - o - 0 0 0 — 4 0 - 0 0 4 4 0 0 i i 0 0 4 0 4 — 0 4 0 0 — 0 0 - 0 - 0 0 0 4 0 — 0 0 0 0 0 0 — — 0 0 — 4 0 0 0 0 0 - 0 4 0 - 4 0 0 0 0 0 0 is 0 4 0 - 4 0 0 0 0 0 19 0 4 0 4 0 0 0 0 0 20 0 4 0 4 0 0 0 0 0 21 0 4 0 4 0 0 0 0 22 0 4 0 4 0 0 0 0 23 0 0 0 4 0 0 0 24 0 0 4 0 0 25 0 0 4 0 0 26 0 0 4 0 0 27 0 0 4 0 0 28 0 0 4 0 0 29 0 4 0 0 30 0 0 4 0 0 31 4 0 0 0 0 32 0 0 0 33 4 0 0 0 indicates that the designated Intensity value falls within the average group - indicates that the designated intensity value falls below the average group 4 Indicates that the designated intensity value falls above the average group 54 quite roughs but it should’ give us some idea of whether this procedure is justified. Examining Table 24 in detail yields some interesting facts, There are no nebulae with all average intensities5 there are 6 nebulae with only one non-zero entry5 and eight more with two non-zero entries (but three of these had a total of only four intensity entries)„ Eleven nebulae could thus be included in the average decrement category r and perhaps two or three others, There were also no completely steep or shallow decrements as described, although several were mostly steep or shallow - or had portions of their decrements that might be so classified. Most of the nebulae have decrements that are mix­ tures of high, average, and low intensities and we do not believe any further attempt to make obvious divisions of the decrements into the steep, average, and shallow would be fruitful.

Aller (1956, p. 129) makes a similar observation in his discussion of planetary nebulae of different degrees of excitement. "The Balmer decrement is closely similar for , . , (all) groups of nebulae, and there is no strong evidence that the high excitation nebulae show a steeper gradient than do the others." Summary and Conclusions

The data Trom each known study containing Balmer decrements in gaseous nebulae have been considered and eval­ uated by virtue of their consistency with other data and by the exactness implicit in each investigator’s general tech­ nique, as well as his care in executing it. Photoelectric^ photographic studies give the best results. Me have concluded that data must be corrected for interstellar differential absorption (reddening) to be meaningful, for such a correction may alter measured- and otherwise eorrected- intensities by more than 100$. Average decrements for forty nebulae were computed using data for those nebular intensities which were corrected for reddening as well as for those which did not need such a correction, i.e., for the high latitude objects. Me deter­ mined that particular groups of decrements do not stand out characteristically from the othersj therefore, we were unable to construct a special subclassification of nebular groups based upon Balmer decrements. Average decrements for the sixteen nebulae studied in detail (group l), the twenty-four others corrected for reddening (group 2) and for those sixteen plus twenty-four in the entire group were computed and discovered to be almost identical (Table 22, page 50). Me thus find nebulae to be objects with very homogeneous Balmer decrements. 56. Graphs of the average decrement for each nebula follow with intensity plotted vs, n; the average decrement for all nebulae from Table 22 also appears in Figure 4, In Figure 5 the logarithmic decrements are plotted for the group 1 and all-nebulae intensity averages. Neither the decrement nor its logarithm is linear with respect to n5 and no other analytic relationship suggests itself to us. 57

200

100 - O

50

20

Oo

20 25 30 n

Figure 5• Logarithmic Plot of Average Nebular Balmer Line Intensities as a Function of n . STARS

Observational Difficulties and Star Types

Balmer emission lines are relatively rare ocearrenees in stellar spectra*, and publication of measured intensity decrements is even rarer. The problems besetting the stellar observer are greater than those for the nebular observer. Relatively simple nebular correctIons, such as correcting lime inten­ sities for the background continuum radiation^ can become quite complicated for a star. In addition to the nebular problem of having other emission lines blended with the Balmer lines5 a stellar Balmer emission line is often split by an absorption line arising from hydrogen, heavier atoms, or even molecules; and these are all superposed upon a continuum. Hydrogen absorption lines are prominent character­ istics for the so-called "early" spectral classes, 0, B, A, and F; Balmer emission lines are very rare. The only meas­ ured decrements for bright Balmer lines in these stars have been for members of the Be class. (The e indicates that emission lines appear superimposed upon the normal absorp­ tion spectrum.)

58 59 An Investigator observes that most stars are con­ stant with respect to all the intrinsic stellar variables he can measure; color, luminosity, radial velocity, size, magnetic field, etc. Some stars, however, are observed to change in one or more of these quantities, and these are called variable stars. The majority of stars observed with

Balmer line emission belong to this group, In the variable stars considered here, the spectrum undergoes continual change, thus introducing uncertainties in any direct comparison of observations made at different times. To give meaningful results, an observer must sys­ tematically record the spectral change as a function of time over some period. This has not always been done. Therefore one is forced to use so-called "average” decrements that arbitrarily lump together intensities measured over intervals of up to several years. Unlike the repeatedly measured and accurately specified nebular decrements# the complicated stellar decrements have not been specified with a high . degree of precision.

We believe that presenting here all of the stellar decrements that have been published would be of little value, and we attempt to restrict the data by using to best advan­ tage the guidelines presented in the previous chapter for data of highest accuracy. Unfortunately, the most signifi­ cant considerations for nebulae observation have almost been 6o ignored for star decrements„ Firsta reddening corrections are virtually non-existent; we indicate in which articles they were considered. Also* photoeleetrieally calibrated data are also rare. Bough photoelectric measurements have been published for two Novae s and some other data exist> but not yet in published form (Aller). #he only result from the nebulae chapter that can be utilized is the realization that recent photographic studies» when carefully made, are considerably more reliable than early ones. For the nebulae5 1950 seemed to be a reasonably good dividing date between acceptable and non-acceptable studies. For that reason, it has generally been used here* A few earlier extensive works are cited. In almost all cases we recognize that this cut-off date is very arbitrary. Most of the studies made after 1950 are probably little better than their predecessors» and the crude eye-estimates found In some post-1950 publications are probably worse than data of some earlier works. In addition to data published prior to 1950, most decrements with less than four Balmer lines have been neglected in this compilation. Data from Be stars are presented here first, followed by decrements from variable stars. Subclasses of variable stars differ considerably from one another. Our categoriza­ tion of these objects into subgroups is principally that 61 used by Kukarkin and Parenago (Strand 5 1963$, chapter 18) whose main division is that between eruptive and pulsating variables. Hot all of their subclasses are represented^ for not all5 perhaps not even mosts variable stars exhibit Balmer line emission. As each subclass is presented in turn,' some of its distinguishing properties are described prior to published decrements for its members.

Be Stars

Be stars for which data is here compiled and the publications used are summarised in Table 2 5*

Table 25. Be Stars Examined and References Thereto

12 Aur Burbidge and Burbidge (1953b) 11 Cam Burbidge and Burbidge (1953b) m CMa Burbidge and Burbidge (1953a) Y Gas Wellman (1952) 56 Eri Burbidge and Burbidge (1953a) V380 Orl Herbig (i960) 48 Per Burbidge and Burbidge f1953a) B Psc Burbidge and Burbidge (1953a) 105 Tau Burbidge and Burbidge (1953a) MD 37998 Burbidge and Burbidge (1953b) HD 58050 Burbidge and Burbidge (1953b)

Wellman’s (1952) data for one object are by far the most extensive offered for an object in this class of stars. Although he observes that "the Balmer decrement scarcely changes," he finds that "the general (absorption) line structure clearly changes," with some absorption line 62 Intensities changing by more than 30^, Either photographic errors are responsible for this finding or y Gas should be classified as a variable star.

Herbig {i960) describes V 380 Gri as a Be star asso­ ciated with nebulosity but Indicates that the intensities he gives are only crude estimates,

In the Burbidges1 (1953a* 195Sb) studies considerable tare was taken. Evidence for reddening was not found. Equivalent widths were measured (see appendix) and in Burbidge and BWbidge (1953b) corrections were made for the superposition of.absorption lines upon: the emission spec*- trum. They found that 11 Gam was barely changed in the tea months between their observations. Their error in decrement for this object is estimated at ±C$> while those for the Other objects in their second study is 30$. Mo estimates

are given in the 1953a study.

Variable Stars

Few pulsating variable stars exhibit hydrogen emis­ sion lines. Mira type stars are one exception to this generalitys but no measurements of their spectra have been found in publications subsequent to 1950. Another exception

is the group of 16 P Canis Majoris type variables5 one exam- . pie of which is included here. Most of the eruptive variable types* on the other hand* evince Balmer emission lines. 63 fa'ble 26 „ Balmer Line Intensities and Decrements For y Oassiopelae from 9/37 to 12/38 Balmer Intensities Computed from Equivalent Widths n (B) (c) (D) (E) (F) 9/25)37 11/28/37 1/22/38 2 /1/38 11/28/38 12/12/ 642 570 823 3 547 4 100 100 100 100 98.2 94.7 82.9 100 100 36.0 24.0 31.5 2 9 .6 26.4 5 37.1 . 35.3 29-5 25.9 6 19.9 16.9 14.6 : 12.4 10,4 23.7 18.1 12.2 10.7 7 10.3 10.8 9.6 1:1 8* 14.9 12,8 " 14.2 ' 2 0 .7

*blended with He 1

Balmer Decrements Computed from Central Intensities ' ' : ' '(A) + (B) (C) 4 (D) (E) 4 (F) 3 313 245 313 4 100 100 100 5 49.3 ,38.3 39.6 6 31.1 20.6 21.3 7 18,1 16.1 17.1 8 12.7* 11.6* 1 0 .5* ^extrapolated from curve of others Table 27= Balmer DeerenBrits of 10 Be Stars

HD HD V380 12Aur 110am 37998 58050 wCMa 56Eri 48Per 9 Psc 105Tau Ori

480 185 202 264 289 500 4 100 100 100 100 100 100 100 100 100 100 65 32.8 42.8 44 51*7 31.2 50 1 38 38 25.5 42.0 12 .9 38 38 31.2 12.2 14 30.7 14.0 88^ I 72 18 31 9 92 24 12 24 10 80 22 23 II 80 28 14 27 12 81 17 11 22 13 22 94 17 8 !! 13 11 25 h 11 8 17 61 8 27 ii 41 15 S 17 19 32 19 10 9 20 31 13 9 21 27 5 22 24 7 23 5 24 7 25 3 Burblclge and Burbldge Herbig '^lend (1953b) (1953a) (I960) Decrements published by Calculated from raw Calculated authors including continuum equivalent width from estimated and absorption".line corrections intensities 65

140

120

80

60

40

4 6 8 10 1 2 1 4 1 6 20 2 4 n Figure 6. Balmer Decrements for Some Be Stars. + 12 Aur □ HD 37998 x y Cass (11/28/37) ° 11 Cam A HD 58050 66 In the literature It is emphasized that there is often difficulty in classifying variable stars into existent subclasses. The classification in this study is done on the basis of explicit statements in the published studies upon which we are reporting.

To be of most value, published data should not only be more accurate, but should explicitly consider stellar variability by giving intensities as functions of time. For periodic and semi-periodic variables the duration of observa­ tions should be a sufficient number of cycles to establish a pattern of intensity variation. In too few studies is time variation considered. Often only a few observations are made; in other cases, time variation is averaged away. Available data, grouped according to variable star subclasses, follows. Pulsating Variables - P Canis Majoris Type

Sixteen of these variables are known (Kukarkin and

Parenago in Strand, 1963). The spectral types range between Bl and B3, and the luminosity periods between 0.15 and 1.25 days. These stars also exhibit the interesting characteris­ tic of having two periods differing by only a few minutes, thus giving rise to a beat frequency in the luminosity curve. Wilson and Seddon (1956) made measurements of the equivalent widths of the first three Balmer lines of P Oephei. These appeared to be constant throughout the cycle. 67 but an intrinsic variation having a long period, perhaps six months# was noted for them-

Table 28. Balmer Decrement of P Oephei

n W ( Standard Error. Integration Range 3 354 ± 37 241 4 100 db 23 141 5 37 ± 2 0 121

Eruptive Variables - T Tauri Stars The T Tauri variables, according to Kukarkln and Parenago, belong to spectral classes B to M and are charac­ terized by irregular light variation from fractions of a magnitude to three or four magnitudes in amplitude. .'The large light variations may occur in a span of a few days to a few weeks. Other characteristics may be found in Kukarkln and Parenago (Strand 1963, page 345) or in Joy (1954). B5hm (1957) gives very low dispersion decrements for VY Orionis and NX Monooeri» Only corrections for atmospheric extinction are made. (See Table 29). Eruptive Variables - Novae Description Novae are stars which flare up suddenly in bright­ ness, perhaps by 10-12 magnitudes, in several days? and which then gradually return to their former state, taking from 68

3 0 0

A

200

100

n Figure 7* The Balmer Decrement of a P Canis Majoris Star • P Cephei 69 Table 29• Balmer Decrements of Two T Taurl Stars NX Monoceri VY Orlonls

1956 n 377" "H7TT 4 100 100 100 100 60 41 51 1 37 31 2? ^ 1 I 3p 26 27 34 9 3^ 26 33 10 19 8 !i II 11 17 11 several hundred days to several decades. The spectral development of these objects is complex. McLaughlin (in ©reenstein 1961) says* The earliest records of spectra have been obtained near the premaximum halt or the equivalent magnitude. All novae observed at this stage have had early-type (B and early A) absorption spectra strongly dis­ placed toward shorter wavelengths . . . All novae near maximum have strong shortward-displaced absorp­ tion spectra of class A or F (with certain peculiar­ ities) without conspicuous emission lines. With the first fading from maximum the spectrum is quickly transformed into the ’typical nova’ spectrum, a pattern of bright lines widened symmetrically about their normal positions and flanked on their sides of shorter wavelength by strong absorption lines . . . A third pattern named 11 diffuse enhanced" and a fourth named the "Orion spectrum" appear. Broad, symmetrically widened emissions accompany both (stages) . . . The hydrogen emissions are the most conspicuous of (those in) the diffuse enhanced system . .".'The spectrum gradually transforms it­ self into one like that of a gaseous nebula, except for the large Doppler widening of its lines. The 70

100 A

50

J L J I I L 10 1 1 n

Figure 8 . Balmer Decrements of Two T Tauri Stars. o NX Monocerci - 3/7/56 A NX Monocercl - 11/11/56 □ NX Monocercl - 11/6/56 x VY Orlonis - 1956 71 transition state is now past and during the.final decline the bright lines weaken more than the con­ tinuum . » . Several novae at the end of decline have relatively narrow bright lines of hydrogen . . . that, persist for many years, while others show only the high temperature continuum. .Study of novae features is made difficult by un- repeatability. No second cycle allows a self-check or chance to average data. A few stars exist which are con- - sidered to be '‘repeating novae" but even for these objects, it is questionable, that subsequent outbursts of such huge proportions would be physically identical to an earlier one. Because of the uniqueness of each nova’s life-, history, it is difficult to try to average Balmer.line intensities for different novae. If one were certain of what stage in the nebular life-span were represented by published spectra, such averaging might be worthwhile. Novae data quality is very heterogeneous. Intensity correc­ tions are most often not present, or are incomplete and inadequate. It is believed that the long-duration, photo­ electric study of Meinel (1963) was an excellent step in the proper direction to secure worthwhile data. Published Data - McLaughlin (1953) made eye-estimates of the Balmer line intensities on plates made from 1919 through 1950. Accuracy is estimated good to one significant digit. Swings and Jose (1952) reproduce intensities measured in 1940, 1942, 1947, and 1949 and add to those their Table 30» Studies Made of Novae Published Since 1950

Nova Studies Made

Nova Aquilae 1918 McLaughlin (1953) Nova DQ Hereulis 1934 Swings and Jose (1952), Bajenov (1956), Greenstein and Kraft (1959)

Nova Hereulis i960 Meinel (1963) Nova DK Lacertae 1950 Wellmann (1951), Larsscm-Leander (1953, 1954) MacRae +43°1 Greenstein (1954) Nova (RT) Serpentis Grandjean (1952) 1909 Nova Scufci 1949 Colacevich (1950) 73 results from seven spectrograms taken during 10 days in May of 1950. From these plates they "integrated over regions of different physical conditions" to arrive at another set of intensities, It is not clear how this was accomplished. Bajenov (1956) remeasured fifteen plates taken between 12/34 and 4/35. ©reenstein and Kraft (1959) took equivalent widths of four Balmer lines photographed at two different phases in the eclipse cycle of the unique binary nova. They remark that great weight cannot be given to the Balmer line strengths found during eclipse by this method alone. There is an "uncertainty of the colors and . . . lack of fiducial continuum under the emission lines." Melnel (1963) took a large number of photoelectric observa­ tions from March through June of i9 6 0 . The instrument used for this purpose was "primitive" and plagued with common photoelectric problems. Observations were made of the first five Balmer linesj atmospheric and instru­ mental corrections were made. Reddening is not men­ tioned . The published results of this study were graphs of the ratios l(Ha)A(HB)s l(Hy + [OIH] )/l(H£) 5 I(H6)A(HP), and l(He)/l(HP) ~as functions of time. From these graphs data are extracted and tabulated here. Wellman (1951) took a series of low-accuracy spectrograms from January to.July, 1950, in which Balmer line intensities were given as a function of nova develop­ mental. stage„ . Larsson-Leander (1953) gave decrementa1 intensities for two days while the nova was in its early post-maximum stage Corrections were made for the continuum and absorption lines. ■ Larsson-Leander (1954), using the same data reduction tech­ niques , gave decrements for subsequent nova stages. Calibration of his data was photoelectric. Greenstein (1954) gave equivalent widths for four Balmer lines in an object whose visual magnitude was declining No corrections were made. These data are included only in the appendix because of the missing H4 data. Grandjean (1952) made eye-estimates from one spectrogram taken long (41 years) after eruption. Colacevieh (1950) made spectrograms over six nights, start­ ing eight days after first signs of the nova spectrum. He presents uncorrected data from very low dispersion spectra. '

The extensive photoelectrically calibrated data of Larsson-Leander shown in Table 33 should be more meaningful than Wellman’s study. 75 Table 31. Balmer Decrements of Nova Serp 1909 and Nova Sen 1949

Nova (RT) Serpentis 1909 Nova Scuti 1949

n 1950 8/6 to 8/14, 1949 3 420 133 4 100 100 5 125 53 6 83 40 7 67 8 50 20 9 25 20 10 8 14 11 8 14 12 8 14 13 8 7 Table 32. Balmer Decrements, of Nova Aqu 1918 n 1919-20 1921 1922 ' 1923 ' 1925 ■11934 '1937 • 1938 1941 1942 '1947 1949 1950

4 (trace) Of 100 100 100 P P 100 100 P 100 100 5 x (trace) Of 200 >100 >200 PP 300 130 P 200 <200 6 Of 75 75 P P 0 67 P 150 <150 7 — Of P 50

p — "present# but not estimated”

-4 o\ Table 33• Balmer Decrements of Nova DK Lac 1950 Wellman n 1/29 2/5 2/17 3/1 3 /1 8 4/8 5/2 6 /3 0 7/3

4 loo loo "™Too“~ T&5 “ TO) TO) "Too" TOTO 5 67 100 100 100 67 67 67 6 67 100 100 67 50 50 53 53 100 7 53 75 67 4o 25 25 27 20 20 8 20 75 53 27 9 20 50 27 7 10 14 25 11 7 5 12 7 13 7 Super­ Early Later n iir™ Commencement of Early Fully developed giant Orion Orion Stage Nebular Spectrum Nebula Nebula Spectrum stage stage stage Spectrum

Larssoh-Leander (These are averaged from many individual values.) Stage n 03 (1/25) Q 3 (1/27) Q4 Q4-5W . Q5i_ Q.6W 3 ~ “T61 183 "TO36™ ^ 0 9^ 4 • 100 100 100 100 100 100 5 47 30 70 73 63 47 6 30 23 56 68 42 7 45 44 25 19 8 30 31 16 11 9 30 - I g U - 6 Early post-maximum

-

n _ 3/10 _ 3 /2 0 3/30 4/10 4/20 4/30 3 220 330 570 : 600 530 .470 4 100 100 100 100 100 100 5+[0III] 29 32 4o 45 52 57 6 14 19 30 4o 42 43 7 9 9 13 18 19 19

I960 n 5 /1 0 5 /2 0 5/30 6 /1 0 6 /2 0 3 400 360 310 280 270 4 100 100 100 100 100 5+CoiII] 62 67 71 74 78 6 44 45 45 46 46 7 19 19 18 17 16

DQ. Herculls 1934 Is an especially Interesting object. It is a close binary and variable star with many unique features (Kukarkin and Parenago, in Strand, 1963). It is apparent that the Balmer line intensities shown in

Table 35 vary considerably in quality. However, due to the complex nature of this nova, it is Intrinsically more diffi­ cult to extract any set of physically meaningful decrements despite use of high accuracy data, , Eruptive Variables - Nova-Like Variables The members of this sub-group vary greatly and, indeed, are classified into.smaller sub-groups. The spectra of the members of this group are similar to the novae, and Table 35- Balmer Decrements of Nova Herculis 1934

"1934™ ^ 1 9 3 5 12/22 1/2 1/6 1/10 2 /1 2/13 2/20 2/22 3/1 3 /2 3/4 3/6 3/15 3/28 4/1 3 "236" 87 74 4 100 100 100 100 100 100 100 100 100 100 100 100 100 100 100 5 75 93 69 74 50 66 65 50 70 83 72 .. 48 6 71 45 115 61 63 30 46 32 32 33 32 31 34 45 31 7 30 .33 71 49 58 23 24 22 25 27 23 30 20 33 17 8 35 9 15 17 15 9 13 12 11 12 9 9 Bajenov” (1956)

— T 9 5 9 ~ ~ 1940 1942 1947 1949 1950 ,0 .3 2 phase .0.04 phase average 3 190* 500* 1000* 4 100 100 100 100 100 100 100 100 5 80 70 70 100 100 106 82 94 6 90* 80 * 80 * 200* 250* 85 46 66 7 90 30 30 80 20 71 50 60 8 80 * 30* 30* 70* 40* 9 40 20 20 50 10 10 30 20 20 30 0 11 30 20 0 12 80 * 30* 30* 50* 20* Swings and Jose (195^7 ™~' dreens^e'in and Kraft (1959) ^blend

3 Figure 9. Typical Balmer Decrements for Some Novae o Nova (RT) Serp 1909■ A Nova DK Lac 1950— 1/29/50 D Nova Scu 1949— 8/6 - \| Nova DK Lac 1950— Stage q.4 8/14/49 x Nova Her 1934— 1959 -Y- Nova Her 1934— 3/6/35 o Nova Her i960— 3/10/60 80

100

50

3 4 5 6 7 8 9 10 11 12 13 n Figure 9. Typical Balmer Decrements for Some Novae 81 some of these variables have sudden flare-ups. Under this heading decrements have been found for members of the following sub-groups: recurrent novae, symbiotic stars, P Cygni stars, and R Coronae Borealis stars. Recurrent novae are very similar to ordinary novae except for the characteristic suggested by the name, Balmer decrements for two of these objects have been found. McLaughlin (1953) gives estimated intensities for four hydrogen lines from 1946 to 1952 in T Goronae Borealis. Bloch and Mao-Lin (1953) summarize equivalent width data from 1946 to 1952 for the same object. They plot in their article the intensity variation of HP, Hy, H6 as func­ tions of time for 45 nights of observation. Griffin and Thackeray (1958) give uncorrected inten­ sity estimates for ten low-dispersion spectra of RS Ophiuchi

1958 photographed between 7/19/58 and 8/29/5 8 , but suspect interstellar absorption because of the steep Balmer decre­ ments . The years of outburst for T Coronae Borealis were

1866 and 1946 and the mean cycle is estimated at 2 9 ,0 0 0 days. RS Ophiuchi outbursts have occurred during 1898, 1933, 1958, resulting in a mean cycle of 11,000 days. R Coronae Borealis Stars differ from the recurrent novae in that they remain at maximum brightness most of the time. Kukarkin and Parenago point out, "Their minima are of Table 3 6 * Balmer Decrements of T Coronae Borealis

“T946 1947 T 9 W ""1950^ m i 1952 n 1946 1947'. .1948" 1949 8 / m 9/17:.r 10/2 7/29 6/2.' 4/16 ■ 5/8 5/4 4 100 100 100 100 100 100 100 100 100 100 100 100 5 34 40 48 44 67 57 89 72 83 67 75 83 6 24 22 28 24 55 36 56 43 83 67 53 33 7+0aII 20 18 , 1 5 17 11 22 25 33 25 0 Bloch! and Mao-kin (1953) ' ~ ~ McLaughlin (1953)

Table 37<■ Balmer Decrements of RS Ophiuchi 1958

- - 1 9 3 8 ------—— ------— ------n 7/19 7/25' 7/28 7/30 8/10 8/11 8 /1 9 8/24 8/29 T ” ” 400 600 “200 50(5 333 300 800 4 100 100 100 100 100 100 ; 100 100 100 5 40 47 50 33 62 67 40 40 33 83

300

A

200

100

3 4 5 6 7 n Figure 10. Representative Balmer Decrements for Two Recurrent Novae. • T Cor Bor - 1946 + RS Oph - 7/19/58 □ T Cor Bor - 10/2/46 x RS Oph - 7/30/58 a T Cor Bor - 5/4/52 o rs Oph - 8/11/58 84 different durations» varying from 1 to 9 mag. For a given star there are times when the minima follow one another rap­ idly, and there have been eases when there was no minimum during ten or more years." Deeremental hydrogen Intensities for only one of these objects was found in the literature. Merrill (1951a, 1961) twice published rough, uncorrected Balmer decrements for xx Ophiuchi. In the first of his studies he found these lines to be "complex, unsymmetrical, and subject to large variation. In 1949, . . . weak; in 1950 . . . wider and stronger." This variability was also noted in 1961.

In Table 38 all published intensities except I(H5) are reported as identical, despite the slight shifting in wavelengths between the two studies, and short term inten­ sity variations. Z Oygoj. Stars are added by Struve and Zebergs (1962) to the nova-like variables category despite their underlying B-type absorption spectra upon which emission lines are super­ posed. Beals (1955) defines these stars by virtue of the presence in their spectra of certain emission line profiles. Any of four line shapes being present is sufficient to clas­ sify an object in this class. These are illustrated in Figure 12. Three objects in this class have been found with pub­ lished deeremental intensities: P Cygni, HD,51585, and X Ophiuchi. Sources of data are again few. Table 38. Balmer Decrements of XX Ophluohl (HD 116114)

observe- n tions 1951 1961 4 100 100 5 50 40 6 25 25 7 18 18 8 14 14 9 12 12 10 9 9 11 7 7 12 6 6 13 6 6 14 5 5 15 4 4 16 4 4 17 4 4 18 4 4 19 4 4 20 3 3 21 3 3 22 2 2 23 2 2 24 2 2 25 2 2 26 2 2 27 1 1 28 1 1 29 1 1 30 1 1 31 trace trace 8 6

100

i

50

0 4 6 8 10 12 14 16 18 20 22 24 26 28 30 n Figure 11. Balmer Decrements of an R Coronae Borealis Star.

O XX Ophiuchi - 1951 x XX Ophiuchi - 1961 37

B

normal emission continuum position level

C D

c ont inuum level

Figure 12. Characteristics of P Cygni Stars’ Emission Line Profiles Type A: Absorption component is on the violet edge; most common type of line in spectrum. Type B: Type A is superimposed on a broad shallow absorp­ tion line. Associated with hydrogen lines only. Type C; The absorption minimum is slightly displaced to the violet. (Dotted line shows case where absorption minimum dips below continuum level.) Type D: Two variations of unusual spectrum lines occur. 88 Burbidge and Burbidge (1955) made limited measurements in April and May of 1953 of both P Cyg and x Oph. Equiva­ lent widths were obtained for emission above the con­ tinuum by considering these stars as radiating black bodies. They corrected for absorption lines and wings. Burbidges cite that evidence is against reddening for X Ophiuchij, whereas for P Cygni there is some dispute.

Kupo (1959) made 79 usable decremental determinations of X Ophiuchi in three months. During this time he dis­ covered that the level of the emission lines above the continuum varied greatly whereas emission line intensi­ ties remained rather constant after he corrected for the superposed absorption lines. The equivalent widths given were normalized tq the unit continuum and then multiplied by the ratio of the measured intensities at corresponding wavelengths. He notes that no absorption

component affected the HP line5 therefore, in distinc­ tion to other lines, its W r(X) was determined from that part of the line which lay above the continuum level. Each of his decrements is thus given as a lower limit. Shapes of the other Balmer lines changed greatly, and his article demonstrated this by tracing the development of the Hy profile (not repeated here).

Arkhipova'(196$) gives low dispersion equivalent widths and a continuum-corrected Balmer decrement for one plate taken in October, 1959s of the P Cyg star HD 51585. 89 Table 39• Balmer Decrements of P Cyg and HD 51585

P Cygni HD 51585 Burbidge and Kupo (1959) Burbidge (1955) 1955 October 1959 3 158 190 4 100 100 5 51 35 6 28 14 7 18 8 9 8.2 10 4.2 11 6.4 12 4.6 13 3.2 14 3.0 15 3-0 16 17 1.4 18 1.9 19 1.1 20 1.4 21 1.1 22 0.8 23 0.3 24 0.3 90

Table 40. Balmer Decrements of X Ophiuchi

ID 24359... ID 24360 Averages 1955 ...62 ...63 ...64 ...80 ...82 ... 86 ...95 ...9 6 ...05 .. .06 ...0 7 .. .08 . . .0< . .15 ...18 ...19 . . .21 . . .22 ...43 ...44 ... 46 . . .47 ...53 #1 #2 3 229 566 4 100 100 100 100 100 100 100 100 100 100 100 100 100 100 100 100 100 100 100 100 100 100 100 100 100 100 5 38 36 31 28 27 30 14 27 27 35 15 21 23 21 40 31 27 35 31 35 28 32 35 27 24 6 29 15 12 21 10 12 7 8 20 32 8 22 17 15 16 9 13 14 10 7 19 15 9 10 5 10 10 20 8 9 9 12 10 7 8 17 9 15 10 15 11 14 12 13 13 15 14 16 15 16 16 14 17 15 18 13 19 10 20 8 21 4 ny Average #1: my average of Kupo's data o g Average #2: Kupo's published average hCCJ '—^ r4 T3 •H bO

3 91

210

180

150

120

90

60

30

3 4 5 6 7 8 9 10 1 1 12 13 14 15 16 17 18 19 20 21 22 23 24 n ------Figure 13- Balmer Decrements of P Cyg Stars

A P Cygni - 1955 o x Ophluchi - 1955 □ HD 51585 - 10/59 4 x Ophluchi - average #1 92 As h%g been-not,ed5 •dqy.lvalent widths lb.:.the:two articles covered in Table -40 were defined differently with Kupo1s being given as a lower limit. Symbiotic stars differ considerably from the other classifications of nova-like variables„ These stars have rather stable light-curve cycles and vary almost continuously. They derive their name from their "combination spectrum"s high-excitation lines superposed upon a low-temperature absorption spectrum (usually of type M). In volume 51 of Das Handbuch der Physikj, Ledoux observed that the hydrogen lines first appear in the cycle while the stars1 brightness increases, reach maximum bril^. liance shortly after maximum light, and then fade out before minimum light is reached. The intensities of the various Balmer lines are cited as being "entirely anomalous": Ha is relatively weak, HP nearly invisible, H6 very strong. He practically absent, etc. As a by-product of citing decrements published since

1950 for these objects, we will be able to comment upon the extent to which Ledoux1s comment is valid. Of the eight stars upon which this section focuses, one (BD-H1°4673) changed from a Be to a combination spectrum during the course of a few decades, and another (AG Persel) made a similar transformation from a P Cyghi star. 93 The stars and the studies made of their decrements are summarized in the following table.

Table 4l. Studies Made of Symbiotic Stars, Published Since 1950

Star Studies Considered Z Andromeda© Merrill-(1950). Bloch and Mao-Lin (1951), Aller (1955) BP Cygni Merrill (1950). Mao-Lin and Bloch (1954), Aller (1955) 01 Oygni Merrill (1950), Mao-Lin and Bloch (1954), Aller (1955) X Cygni Mao-Lin (1950b), Fugita (1954) A0 Pegasi Mao-Lin (1950&), Mao-Lin and Bloch (1952), Burbidge and Burbidge (1954), Arkhipova and Dokmchaeva (1962) AX Persei Mao-Lin and Bloch (1954, 1957) MW0 603 Tifft and G-reenstein (1958) E M I 0 4673 Merrill (1951)

The quality of data given in these thirteen articles varies greatly. The articles are considered chronologically,

Merrill (1950) gave crude, uncorrected, average intensities for three stars observed on as many as eight occasions during a span of three years. Although only the aver­ age data are given, he remarks, "The relative intensi­ ties . « . are subject to considerable variation from time to time» Moreover, estimates . . . are subject to considerable variation ..." Mao-Lin (1950b) estimated the Balmer line intensities in his study of the post-(1948) maximum spectrum of x Cygni. Mao-Lin (1950a) similarly estimated intensities from the heights of microphotometer tracings for AG Pegasi„ No corrections were made for either study. Bloch and Mao-Lin (1951) measured central intensities above the continuum after making only a correction for their grating dispersion. Decrements are given for 1946 (average of three plates in as many months) and 1948 (one plate). Merrill (1951) gave his estimation of the average decrement for an eight year span for BDfll* 4673 * Merrill has followed the development of this unusual star for about

30 yearss but his published intensity data are of little use. Mao-Lin and Bloch (1952) made a study for'AG Pegasi compara­ ble to their one in 1951. They note the evolution of its spectrum in which different lines appear in differ­ ent years. This was a low dispersion5 unconnected study. Fugita (1954) presents the changes in the intensity decre­ ment of x Gygni during two months in 1950 near the time of its light maximum. His five sets of data are unconnected and are given in units of the surrounding continuum. Mao-Lin and Bloch (1954) presented decrements for three more stars, measured from plates taken during October of 1952. r - 95 Burbidge and Bnrbidge (1954) took spectrograms of AG Pegasl and measured the equivalent widths of the BaInter series from H4 to H30<> By then assuming that the continuous background arose from a black body at 5 0 ,0 0 0 K, they obtained their "observed Balmer decrement." They noted in their study the trend toward higher excitation over the years and also stated that they believed the self­ absorption effect in this star was of minor importance * Aller (1955) presents data obtained over several years for three objects. Equivalent widths are generally given for the highest dispersion data, but most of the num­ bers presented are intensity measurements from medium and low dispersion spectrograms. Some of these are eye- estimates. Averages are not given but rather the individual dated measurements. Mad-Lin and Bloch (1957) used microphotometer tracings to find four sets of uncorreeted intensities presented for AX Persei from 1952 - 1957« Tifft and Greenstein (1958) presented two decrements for this star, whose spectral type they place between the high- excitation symbiotic spectra of AX Persei and Cl Cygni and the lower-exe it at ion object RW Hydrae, A reddening correction was considered and discarded, since it led to discrepancies with the observations„ Arkhipova and Dokuchaeva'(196$) gave observed equivalent width Balmer line decrements for one night in each of 96 the three years 1956, 1959 (two plates), and i9 6 0. Their instruments were low dispersion and they dis­ covered an unexplained 20$ equivalent width deviation for the line measurements taken in a single night, 8/27/59. Their large measurement errors are not corrected for.

Decrements published for these symbiotic stars follow. The Mao-Lin and Bloch data are in good agreement ' with those* of-Merrill' (see Table:. 43) • Aller ’ s results are self-consistent. It is not clear if the discrepancy is due to real time variations or photometric inconsistencies. Ledoux’s allegation of completely anomalous Balraer decrements for the symbiotic stars is not seen to be generally borne out by the data which we have found. Although Pugita1s study of x Cyg gives support to Ledoux8 theory, most of the decrements of these stars seem merely to be shallow (with respect to the gaseous nebulae), and their intensities do generally decline with increasing n= One suggested mechanism for producing such decrements has been self-absorption of the early Balmer lines. De-excitation mechanisms have also been mentioned. We know of no unique mechanism that poses a generally accepted solution for these decrements and we strongly doubt such a mechanism coming to Table 42. Balmer Decrements of Z Andromedae C

MerrilT Bloch Alier (1950) Mao-Lin [M5(X)]. 1951 Average 1946 1948 9/24/50 9/10 9/13 • 3 500 327 4 100 100 100 100 100 100 5 50 4? 63 44.2 50 32 6 33 22 22 14.3 2 7 .7 2 5 . 7 25 14 13 35 21 8 20 18 11 15 9 17 13 10 17 6.3 11 15 6.3 12 13 6.0 13 13 7.8 14 12 4.3 15 12 ' 5.3 16 10 3.7 17 10 0.1 18 8 0.1 19 8 0.1 20 8 . 0.1 21 8 0.1 22 6 23 - 6 24 5 25 5 26 3 27 3 28 2 29 2 30 tr 31 tr •<2 Table 42» (Continued)

T9ffT ...... _ _ . „ 195T"' 7/27 9/22 77^10/3 7/30 10/3 10/1% 2/3 g/t 7/24 3/29 5/9 f7^9 3 4 100 100 100 100 100 100 100 100 100 100 100 100 100 100 100 100 100 5 75 75 50 47 30 29 5 0 40 40 50 40 35 20 50 37 33 35 6 20 15 15 10 15 6 10 8 17 15 7 10 5 5 4 3 0 .3 I 4 5 8 5 2 3 2 1 (low dispersion eye-estimates of intensity)

VO oo 99 Table 43. Balraer Decrements of BP Gygnl

HerriIT Mao-lLin . Bloch 91 M 2 . ~57gi 677 T7ir~87rr 3 500 213 4 100 100 100 ioo 100 100 100 5 50 26 33 29 29 33 51 6 30 11 112 20 17 16 29 7 20 15 5 6 15 12 18 8 17 3 3 3 10 4 14 9 13 i- i 3 6 10 13 11 11 12 11 13 10 14 10 15- 9 16 9 17- 9 18 7 19 7 20 7

21 “ 23 5 24 3 25- 28 2 Table 44. Balmer Table 45. Balmer Decrements of Cl Cygnl Decrements of M-JC 603 Merrill Aller Mao-Lin Bloch 1948- 1950 T ^ L — aftT- " 7/16 8/12 1952 Greenstein 1949 9/19 October 3 500 271 4 100 100 100 100 100 100 4 100 100 5 50 20 50 50 60 48 82 80 6 33 7 35 60 24 29 1 68 68 7 25 12 20 10 16 8 20 5 6 2 . 8 % 9 12 9 41 10 12 .10 36 11 11 10 11 32 20 12 8 12 27 16 7 23 12 II 7 it 12 7 I? 12 II 7 18 8 .5 14 8 II >5 11 14 8 19 5 19 14 20 5 20 14 21 3 21 14 22 3 22- 3 I! . 3 11- 25-28 2 27 29-30 trace

^•equivalent width 100 teye estimated intensities Table 47 . Balraer Decrements of X C y g a i of M Perse1

Mao^Lxn 1952 1955 11/24/56- 10/10 0/6-10/8 10/23-11/17 1/19/57 . 3, 125 3 4 100 100 100 100 100 100 4 100 100 100 100 5 So 96 90 165 524 174 5 45 48 = 64 58 6 60 234 225 194 1115 198 6 28 31 : 56 43 7 weak weak.weak 3 7 13 v 16 48 ■ 28 8 44 377 277 512 87 8 10 10 ' 67 30 9 40 120 200 145 9 7 9 43 18 10 32 32 85 10 6 36 12 11 36 23 51 11 4 39 10 12 20 12 29 9 13 8 13 2? 7 14 14 35 9 15 16 15 30 7 16 16 7 17 17 II 4 18 8 18 13 4 19 4 19 11 1 20 20 13 3 21 4 21 14 1 99 10 2 102 Table 48. Balmer Decrements ©f AG Pegasi

Burbidge Mao-Lin. Mao^Lln and and Arkhipov and (1950) ' Bloeh; (195^)' Burbidge Bokuchaeva (1962)* "i * n 1946-1948 1946 1948 1951 1954 1956 1959% 1959$, I960

107 182 187 213 100 100 100 100 100 100 100 100 88 61 63 71 63 60 64 92 51 I 81 39 39 50 54 37 57 51 37 75 25 26 46 36 9.3 29 30 25 69 30 31 43 43 9.5 19 19 9 44 21 22 29 17 3.3 10 11 10 50 14 3.8 . 11 11 38 13 4.2 11 12 31 8 4 13 25 6 .3 14 ■ 22 4 19 5 ;! il 15 3 IT 12 2.:1 18 19 2 ,6 19 2.4 20 1.9 21 2.1 22 . 1.7 23 1.5 24 1.4 25 1.3 26 1.3 27 0.7 28 0.9 29 0.9 30 0.8

^equivalent width decrement Table 49. Balmer Decrements of ED -5-11° 4673

Merrill (1951) 1942- I.950 20Qn lOOn 5 60n 6 40n 7 3 On 8 3 On . 9 20n 10 2Gn 11 l6n 12 l6n 13 l4n 14 l4n 15 l4n 16 12n 17 12n 18 12n 19 12n 20 10 21 10 22 10 23 10 24 8 25 8 26 8 27 8 28 8 29 8 30 6 31 6 32-33 4 34-36 2 37-38 trace n = Diffuse line Figure 14. Representative Balmer Decrements of Some Symbiotic Stars (A).

a . z And - Merrill's average D Z And - 9/13/51 A BB 4- 11° 4673 - 1942-50 + AG Peg - 1954 n m Peg - 1946-48 o m e 603 - 1958 104

— ii

’,00

50

0 4 6 8 10 12 14 16 18 20 26 28 30 32 34 36 n ---- Figure 14. Representative Baliner Decrements of Some Symbiotic Stars (A). 105

150

100

50

H 6 8 1 0 I 2 1 4 16 18 20 22 24 26 n Figure 15. Representative Balmer Decrements of Some Symbiotic Stars (B). A BF Cyg - 1948-49 4- AX Per - 1958 O % Cyg - 10/10/53 x AX Per - 1958-7 O AX Per - 1959 . 106 the fore until both observations and calculations become more accurate.

Summary and Evaluation of the Published Data

For the symbiotic variables and for almost all of the stars considered in this chapter* we do not believe that any published decrement may unequivocally be termed defini­ tive. By the same token we do not feel that average decre­ ments for stellar groups would be meaningful at this time. Table 43 for the symbiotic star BF Cygni gives us an example of this later point and helps convince us of the arbitrari­ ness involved in weighting true decremental variations against photometric inconsistencies. The range of inten­ sities for each stellar group is summarised in Table 50. Criticism against individual decrements may be justified on several accounts. Most investigators of stellar spectra have without justification used central intensity measurements to give decrements. Me explained the potential inaccuracies inherent in this process in the introductory remarks to this thesis; we can punctuate those remarks by considering the two decrements published by Aller for V.' 01 Cygni which were measured from one spectrogram (see

Table 45* p. 100 for data measured for 7/16/5 0). One decre­ ment is from the equivalent width ratios; the other is from the intensity ratios. They differ greatly. Equivalent 107 Table 50. Balmer Emission Line Intensity Value Spreads Published for Eight Star Categories The number of stars included in each group is indicated in parentheses

Star Type

n Be (11) 6 Canis T Tauri Novae Majoris (1) (2) (7) 3 185-823 354 & 37 74-636 4 100 100 * 23 100 100 5 24-65 37 ± 20 41-60 29-125 6 10.4-42.8 27-41 1 4-1 0 0 7 7.5-38 26-34 9-90 8 10.5-72 26-34 9-75 9 12-92 8-19 .6-50 10 22-80 11-17 8 -3 0 11 14-80 5-30 12 11-81 7-14 13 22-96 7-8 14 8-67 15 H -8 3 16 8-48 17 9-61 18 8-41 19 9-32 20 9-31 21 7-27 22 7-24 23 5 24 7 25 3 108 Table 50. Balmer Emission Line Intensity Value Spreads Published for Eight Star Categories (cont.) The number of stars included in each group is indicated in parentheses

Star Type n Coronae P Cygni Recurrent Borealis Symbiotic Novae (2) (3) (8) ...... (XL ...... _ 1 " 3 '.... 200-600 1 58 -566 107-500 4 100 100 100 100 5 33-89 40-50 14-51 2 0 -5 2 4 6 22-83 25 7-29 6-1115 7 0-33 18 8-20 0 ,03-75 8 14 17 1-512 9 12 8 .2 -1 5 |-200 10 9 4.2-15 3.8-85 11 7 6.4-14 4-51 12 6 4.6-13 4-31 13 6 3.2-15 6.3 -2 7 14 5 3.0-16 4.3-35 15 4 3.0-16 5.3-30 16 4 14 3.5-23 .17 4 1.4-15 0.1-16 18 4 1.9-13 0.1-19 19 4 1.1-10 0.1-14 20 3 1.4-8 0.1-14 21 3 1.1-4 0.1-14 22 2 0.8 1.7-10 23 2 0.3 1.5-10 24 2 0.3 1.4-9 25 2 1.3-9 26 2 < 1.3-8 27 1 0.7-8 28 1 0.9-8 29 1 0.9-8 30 1 0.8-6 31 trace trace-6 32 4 33 4 34 2 35 2 36 2 37 trace 38 trace 109 width decrements are also given for other stars, but these are incommensurate with the central intensity decrements published on other occasions, e.g. Z And and AG Peg. It is obvious that raw data must be carefully obtained and carefully developed to give accurate decrements. Suspicion must be raised, for example, regarding the accuracy of Arkhipova and Dokuchaeva *s (1962) 20$ line intensity variations in one night which then gave them a 50$ variation in their normalized decrements. While not as apparently incongruous as this case, other published results from other investigators may likewise be expected to contain large inaccuracies. The complexity of the stellar spectra contributes to uncertainties in the data. For example, emission lines over- lying absorption lines demand a high degree of sophistication in data reduction, and only a few researchers, among them the Burbidges, show a willingness to meet this problem directly. Spectral resolution difficulties are also accen­ tuated by complex spectra. The significance of any decrement is impared when other emission lines are irresolvably blended with those of the Balmer series. In some publica­ tions such blends are indicated, in others they are not. When one considers that resolutions vary extensively from study to study (for the symbiotic stars, from 10-300 A/mm), it is necessary to indicate exactly which Balmer lines are blended in each set of data. 110 Poor focus in one’s apparatus can further invalidate central.intensity measurements. Very^rarely in stellar decrement studies was inter­ stellar reddening considered. Some stellar investigations should be mentioned here as positive examples of careful procedures in observa­ tional technique and analysis. We feel the most useful studies which we found were Meinel's and Larsson-Leander1s novae studies» the Burbidges’ work with Be and other starss and Kupo's x Ophiuchi study. Certainly Aller’s yet to be published photoelectric studies will join this group.

Additional Data

Additional studies of Balmer decrements for the stars discussed in this chapter have been made. A list of earlier references dealing with these stars follows and concludes this chapter. m Table 51° Studies of Stars Published Prior to 1950

Be Stars 11 Cam: Karpov (1933) y Gass: Struve and Swings (1932), Karpov (1933), Baldwin (193§) ■ Novae Nova DQ Bereulls 193^:. M a m s and Joy (1936)..# Swings, and Struve ,(19^2a) Swings and Jose (19^9) Nova (BT) Serpentis 1909: Swings and Struve (1942b) Nova°Like Variables I O^gii St^s P Cyg: Blvey ( 1 9 2 8 ) Struve (1935), Struve and Roach (1939), Swings and Struve (1940), mcaaika (1949), Beals (1955) Stars Z And: Swings and Struve (1940, 1941a, 1942a, 1943), Merrill (1947a), Mao-Lin (1949) BP Cyg: Merrill (1943) Cl Cyg: Swings and Struve (1940) X Cyg: Merrill and Burwell (1930), Merrill (1941, 1947b) AX Per: Swings and Struve (1941b, 1942a) ED *11°4673: Merrill (1929, 1932), Swings and Struve (1940) SOLAR PHENOMENA

The closest source of extra-terrestrial hydrogen emission spectra Is the sun. Its relative proximity might mislead one to believe measurement of the Balmer decrement simplef and to suppose that many accurate measurements exist. Self-compatability of existing data is examined, and an evaluation of these data is attempted in this chapter.

Solar Regions and the Sites of Hydrogen Emission Spectra

Only the small fraction of the sun which makes up its atmosphere is directly observable. The lowest layer of the solar atmosphere surrounding the solar core into which we cannot visually penetrate, is termed the photosphere, and its projections on a plane perpendicular to the line of sight from an earthbound observer to the center of the sun is the disk. At the end of the disk, the extreme solar limb, the intensity of the sun decreases by a factor of e in 100 kilometers, which gives the photosphere what appears to be a very sharp edge. Beyond the limb for some 10,000 km the intensity decreases much more slowly. This region, the chromosphere, appears uniform at its base, but as one looks farther out along the solar radius it gradually appears to bristle, porcupine-fashion, with bright spike-like columns, 112 113 er spicules. These spicules penetrate into the edge of the upper chromosphere* beyond which» extending as a faint halo, is situated the third solar atmospheric region, the corona. The coronal intensity fades much more slowly than the chro­ mosphere's and has been observed in eclipse to extend for several solar radii. Smith and Smith (1963, p. 9) suggest that there is some evidence that it extends beyond the earth. Although hydrogen comprises the major portion in the sun’s make-up, the opaque photosphere is such that it ex­ hibits, superimposed on its bright continuous spectrum, strong Balraer absorption lines, and most frequently these mask the ubiquitous Balmer emission of the chromosphere for an observer. Athay (Thomas and Athay 1961, p. 25) points out that the first observation of the chromospheric spectrum was accomplished by Young in 1883, and also that In 1909, Hale and Adams were able to photograph it from a 60-foot- tower at Mount Wilson without the benefit of the moon obscuring the photosphere, Because the spectrum of this portion of the sun is most easily observed during the few moments afforded by a total eclipse it has long been termed the flash spectrum. The boundaries of the duration of a total eclipse are the second contact and third contact of the solar and lunar limbs illustrated in Figure 16. The coronal glow does not have a hydrogen spectrum among its characteristics. 114

Figure 16. Schematic Representations of the Contacts of the Lunar and Solar Limbs during Total Solar Eclipse. A. First Contact B. Second Contact C. Third Contact D. Fourth Contact During second and third contacts only the corona and part of the chromosphere are observable. 115 Hydrogen emission is also observed in various "disturbed" portions of the sun. These may take the form of prominences (quiescent or active) or flares. Most often these will be seen occurring in the chromosphere on the solar limb, but observations are made on the disk of both prominences (which appear there as dark filaments) and very bright flares. Prominences vary greatly from one another. Hoyle (1949) spent several pages in trying to distinguish characteristics of one type from another. All flares, however, involve the extension of matter from the lower chromosphere into the higher chromosphere and even into cooler regions of the corona. Flares differ in that they are bursts of light rather than matter, although, as Kiepenheuer (in Kuiper 1953s chap. 6) points out, matter in the surrounding neighborhood of flares often is ejected. In this chapter observations of the solar chromo­ sphere, flares, and prominences are considered.

The Flash Spectrum of the Chromosphere

History and Observation Techniques Although Hale and Adams observed the chromospheric spectrum outside the eclipse in 1904, we have not found any extra-eclipse observations that have resulted in publication of Balmer line intensities. One reason for this is the observer's predicament in knowing h, i»e., how far beyond the solar, limb he is observing, ©ccultation has been necessary to determine this. We will later show that both the intensities and the decrement vary with height above the limb. An additional problem for extra-eclipse observations is the unavoidable scattered light from the photosphere con­ taining Fraunhofer absorption lines. The occurrence of a total eclipse has never in itself guaranteed reliable data, even after extended preparations, long journeys to optimal observation sites, and perfect rehearsal. Expeditions have failed for a large variety of reasons. All of the data .known to have been published are included in this chapter, but unfortunately the amount of data is meager, as may be seen in Table 52. Two references in the table.are made for "excited" (or active) regions. According to Thomas and Athay (1961), emission in these chromospheric regions extended to unusually great heights and lines of high excitation potential were abnormally strong at all heights. Conversely, lines of low excitation potential were weak. Table 52 lists studies in which we found data. Because quality of measurements is often directly contingent upon the observational technique used to obtain them, a designation of techniques used is presented in column 2 by the following abbreviations: Table 52. Balmer Line Studies Resulting from Solar Eclipse Observations

Normal Regions

Method of Balmer Lines Heights E(0) . ^n Eclipse Viewing Included Measured Inel, Inel» Reference

I905 ca H5-H23 K. Schv/arzschiId (19O6 ) 1926 ea H3-H7 Russell (1926) 1926 as. H6-H36 - C . R. Davidson and P. J . M. Stratton (1 9 2 6) 1905#1925# S. A. Mitchell (1930) 1926 ca H3-H37 1932 jf H3-H31 670-4000km - X 1G. G. 011116 and D. H Menzel (1936) 1936 ns' H3-H26 - A. D. Thackery (1937) 1936 ea H3-H27 290-5950 X 2R. G. Athay, D. H. Menzel, P. Orrall (1957) 1905,1925# ea ■' - ; S. A. Mitchell (1947) 1930,1937 H3-H37 1905,1925# — K. Araki (1952) 1930,1937# H3-H37 V. P. Vyazanitsyn 1941 ea H3-H18 100-7300 (1951) 1945 ea H3-H12 650-6200 X V. P. Vyazanitsyn (1952) 1945 # H4-H16 500-5500 X X % . Kristenson (1955) 1952 jf H3-H31 530-6300 X ^R. G. Athay, D. E. Billings, J. Evens# ¥. 0. Roberts (1954) 1952 ns H4-H11 I2 5-I638 J. Houtgast (1957) 1954 ns, jf . H4-H17 7 3 0-3290 X D. Ya. Martynov and V. Ya. Abluseva (1962) 1954 ns, jf H3-H8 (-1000)-8000 x J. Houtgast (1962) 1961 jf H4-H6 A. L. Stolov (1962) Excited Regions 1932 H3-H31 900-2750 G. G. Clllid and D. "" H. Menzel (1936 and 1937) 1952 H3-H28 290-4510 x R. G. Athay and R. M. Thomas (1957) Region aH "Reworks the data of Hemmendlnger5 1939# Thesis* Princeton U. Considers also Menzel5

^Individual points presented only graphically, ^Data also Included in R. H. Thomas and R. G. Athay, I9 6I, Physics of the Solar 5, Interseienee Piabl, Inc. M. Y. 1 1 8 ca - chromospheric arc mp - moving plate jf ~ Jumping film ns - narrow slit These methods are described and evaluated In the following paragraphs which comment In general upon observational techniques, Slltless spectra comprise the vast majority of spectrographlc data for the observed eclipse flashes, for several reasons, First, the narrow chromospheric arc serves as a distant source and no slits are needed, A series of mono­ chromatic crescents result from the Integrated light strik­ ing a grating, and these are recorded on various portions of a photographic plate, These are illustrated schematically in Figure 17, as are the circular Images of the coronal halo. Since the corona and chromosphere differ significantly physically, the rings and crescents do not necessarily coincide. The corona also gives rise to a weaker continuum which is not illustrated in Figure 17* Secondly, from a slltless spectrum one can determine relative heights in the chromosphere reached by each line by measuring the thickness of each monochromatic slice. Thirdly, it is technically difficult to use a slit because of the difficulties involved with achieving and maintaining alignment. Figure 17o Schematic Drawing Showing the Spectral Resolution of the Solar Image During Eclipse Totality, During totality: Moon is blotting out photosphere, A slice of chromosphere and the corona shine on the grating and are diffracted. The circular images on the plate are due to the corona# crescent images to the chromosphere. 119

velocity of moon = 310 km/sec

earth atmospheric layer

photographic plate

Figure 17. Schematic Drawing Showing the Spectral Resolution of the Solar Image During Eclipse Totality. 120 The earliest flash measurements were obtained from slitless spectra. As discrepancies became apparent improve­ ments were made. For example# large variations were noticed in line intensities from early flash spectrum data. Olllid and Menzel (1936) realized that these spectra might have arisen at different heights above the photosphere» The rela­ tive velocity of the moon crossing the sun permits the cover­ ing or uncovering of 1000 km (or about l/lO) of the ' chromosphere in three seconds s and it soon became apparent that observed intensities depended upon which portion of the chromosphere was being observed. 011114 and Menzel no longer were content in obtain­ ing a single exposure on a plate and then trying to change plates as rapidly as possibles a technique which is simply called the chromospheric arc technique, for purpose of Table 1. They utilized a jumping film technique, similar to a motion picture, in which the film was advanced while the shutter was closed. In this way they and their more success­ ful successors have been able to obtain extended series of exposures in which the moon obscures different portions of the chromosphere, and allows observation of radiation from atoms in distinct portions of the chromosphere. By this technique changes in intensity are relatively easily observed. Another adaptation of the chromo spheric arc technique used a moving plate under a long slit placed above.the plate, 121 By having the slit oriented parallel to the direction of dispersion it was possible to record a continuous record of the change in the spectrum at one region on the limb. The primary difficulty with this method is the impossibility of knowing just how representative is the single region being observed. For this reason few observations have been made in this mode. The limitation on one's ability to change plates rapidly, together with the need to have as many exposures as possible has made the moving film the most popular of these methods. Spectra at many heights in the chromosphere are necessary not only to determine intensity variation with height h above the photosphere, but also because of the difficulty in establishing the h = 0 level. Exposures taken just before and after the appearance and disappearance of the Fraunhofer spectrum of the disk serve to give a great deal of otherwise unattainable accuracy to the zero of the height scale. All solar flash spectra suffer from common faults due to averaging effects. For example, all the emission which we observe on earth is averaged along the line of sight over the varying thickness of the chromospheric layer through which the observer looks. Further averaging of intensity occurs during the tracing of a spectrum by a micro­ photometer due. to the finite slit width of that instrument. 122 The slit less spectra suffer an additional problem. As may­ be easily seen from Figure 17# p. 11 9# without a slit all of the light above the lunar limb is projected onto the plate. The only way to determine which radiation originates at a given ehromospherle height is to subtract from the measured Intensity at height hs the intensity due to radiation occur-? ring above that height. It is easier to screen out all but that part of the chromosphere by a ’’narrow slit" set between the sun and the grating and oriented in a direction perpen­ dicular to the dispersion'. The resulting line profiles as traced by the microphotometer are much more accurate. The use of such a slit is difficult (see Davidson and Stratton 1926)# and only recently has it again been utilised (J.

Houtgast 1957# 1962)# combining its advantages with those of the jumping film technique. - With slit spectra integration occurs due to blurring effects. This is due both to "seeing”* that is, atmospheric scintillation, and to poor focusing.

Specification of ©aimer Decrements in the Chromosphere

It has already been stated that relative Baliner line intensities in the chromosphere vary with their position beyond the solar limb.. The sun thus poses the problem of observing and specifying spatial variations in the Balmer decrement for an extended source. A number of questions are 123 considered in light of the data which follow. A brief state­ ment and discussion of these questions hopefully leads us.to a meaningful way to present the data, after which we examine the data themselves. For each particular eclipse one wishes to know the Balmer decrement at every point in the chromosphere, One should know the intensity contribution of spicules compared with surrounding interspieular regions. The line intensities vary with height5 does the decrement? Is the decrement isotropic, or does it depend on angle? If the decrement is specified spatially for one eclipse, will these observations be valid for the next— i.e., is the decrement a function of time? Are corrections other than instrumental ones neces-' sary for eclipse data, and are intensity specifications well standardized? The methods used to calibrate eclipse intensities vary considerably in quality. It has proved possible, how­ ever, to use two different standardization techniques on one set of slitless data to ascertain a reasonable degree of reliability for each set of measurements (Athay, Menzel, and Orrall 1957). Standardization resulting from impressing the well-known spectrum of a standard tungsten lamp on plates or film has correlated well with a standardization technique in which the dispersed coronal continuum image ±. 90° to the line 124 ©f dispersion was eonsidered to Mire the same energy distri­ bution as the tmeelipsed sun. E^(h)» the coronal brightness at wavelength X as a function of M, was determined absolutely by an external standard* and was then used to calibrate the chromosphere spectrum» Other standardization cross-checks have been made which allow the reader to be relatively con­ fident of data calibrated in these ways. Me express this confidence later::-!® eur weighting of the results from different eclipse observers» Discrepancies exist on the order of 100-1000 km. in determinations of'the zero height‘level. The h = 0 point is set arbitrarily (Thomas and_Athay"Igiol* p. 35) at the point Where the optical depth for a tangential monochromatic ray ©f wavelength 4700 through the solar atmosphere equals unity. Optical depth is defined as being the distance traveled by a /beam of light through , some medium in which the intensity of the beam is reduced to 1/e of its original value. It seems desirable that future calibrat1ons be made this way (or in some other standard way)* for the present literature is not uniform on this point* leaving one calibration method unrelated quantitatively to another. ' An error of 300 km in the zero height can mean an error of between 2 and 3 in the ratio between the intensity of E4 -and the higher members of the Balmer series. ' Such an error is unnecessarily added to the errors that one is ratable to avoid or correct* since 125 this standardization of the calibration is often simply a matter of definition. The possibility of being reasonably correct in establishing heights did not exist for most early observers and for many present-day ones, due to their not obtaining measurements at moments near or during second and third contacts of lunar and solar limbs, the only moments in which the edge of the photosphere may be precisely determined. 011114 and Menzel (1936) and later observers have found that intensity in the spectral lines varies as a func­ tion of height approximately in the following ways

•1(h) = l(0)e~^h .

This relation is found to hold rather well above a certain height, below which it serves as lower limit to the intensity (see for example Kristensen 1955)• Mathematically,.one can . choose either to say £ = 3(h) or to say that the above formula is the first term of a series expansion 1(h) = El^e”^ 1 , where the (6^3 and the (I.} are sets of constants. No one has done more than to consider the first order repre­ sentation and deviations from It $ so our attention is focused upon determining to what extent it is valid, and in finding, if possible, a numerical representation in this form which gives the observed values of the chromospherie intensities. 126 Taking the logarithm of the intensity gives

loglol(h) = log10l(0) - ^2^. h •

Plotting log I vs. h gives a straight line with log10l(0) as the h-axis intercept and -0/2.3 as the slope, ^ ^ • The question of finding the Balmer decrement is immediately complicated. If the above representation is correct at height h,

5l(h ) - In (h) _ In (° ) n m-

Pour constants and the height in the chromosphere must be known to give the decremental relationship at each arbitrary height. Various regions on the limbs must be examined during each eclipse in order to determine if chromospheric inten­ sities are also functions of the azimuthal angle. Athay and Thomas (1961) revised their (1957) conclusion and decided that chromospheric radiation is generally isotropic, and that ’’active” or "excited” regions were exceptional. Athay, Menzel, and Orrall's recalibration of Hemmendinger*s (1939) measurements of 14 regions on the solar limb showed 9 regions of normal radiation. The other five were active and showed prominence features. Similarly, Athay and Thomas (1957) found two of the three regions that they examined to be nearly identical. They cite a study of Dunn's in Hot outside ; . - , ' . is? eclipse as further evidence of symmetry in the chromosphere. Thus5 in our calculations we consider flash spectrum measure­ ments for individual regions as representative„ without

s 1 • ■ . _ concern for limb position if the measurements are not taken in excited regions. Intensity measurements have not been made for the fine structure of the middle and upper chromosphere„ i.e.„ the spicules. Observers have not yet overcome the technical difficulties in resolving one spicule from its neighbors and in gathering adequate light from it during the onrush of the moon’s limb. The question of whether temporal change occurs in t h the intensities has not been decided. It is possible that temporal changes occur over the 11-year period of the sun- . spot';cycle„ but insufficient data exist to determine this. The magnitude of an error in h = 0 of the order ^previously mentioned would most likely obscure any temporal change. Figures 18 and 19 (pp. 149 and 150) are drawn for some of the decrements computed from measurements made for eclipses . ' in 1941 and 1952 respectively. The difficulty in comparison is immediately apparent from the difference in quality of the data. The Published Data

In accordance with the preceding paragraphs the data will be presented without searching for chronological varia­ tions or variations depending on limb position. So-called excited regions will be considered separately. Mon- logarithmie decrements are presented for each publication. We have normalized each decrement by setting Ig^(h) = 100. The original data will be found in Appendix D. Abbreviations will be used on occasion to refer to the sources of data. The initials of the authors1 last names will suffice except when this would lead to ambiguity. In that case the year in which the eclipse was observed follows the initial in parentheses. The data for each publi­ cation will be presented in the order given by the follow­ ing notes.

S - 1905 eclipse. This attempt by K. Schwartzschild

(1906) is the first careful attempt to measure ehromospheric intensities that appears in the litera­ ture . The data were taken at a single height on one plate. DS - 1926 eclipse. Davidson and Stratton (1926) published their measurements and repeated those of Russell (given in parentheses. Table 53* P» 135). It was necessary for us to normalize DS intensities to 129 Russell’s at H6 to obtain this early extended decre­ ment * No designation of the height of observation is made„

M(1926) - 19.05, 1925, 1926 eclipses. Mitchell (1930) by visual estimate of intensity* averages 11 a fair number of plates'' from these expeditions to arrive at this average. He also includes his estimate of the maximum ohromospherie height which each line reaches. M(1947) - 1905, 1925, 1930* 1937 eclipses. Mitchell’s (1 9 4 7) data are not used directly here* but instead we give Araki’s (1952) corrections of Mitchell’s intensities. Araki used R. Wildt’s 19%7 empirical color corrections on the log of the intensity to arrive at his figures. Mitchell’s work is histori­ cally interesting. He had been on ten eclipse expeditions at this time and had made no measurements to try to establish a height dependence for inten­ sity. He also strongly defended the method of eye- estimating intensities; "It is the considered opinion of the present writer that estimates by a trained observer have as high* or even higher* accuracy than that measured by a microphotometer of any kind." Fortunately 01111# and Menzel (1936) and subsequent observers were more aware of oscular limitations. ~ 130 T ~ 1936 eclipse. Because of poor seeing conditions Thackery (1937) was able to salvage only one of several exposed plates and a moving plate spectro­ gram. He does not try to over-analyse his data and the values given here result from our reading a graph in the publication. The height is unspecified. CM - 1932 eclipse. CillitS and Menzel (1936) published the first data in which the Balmer line intensities were measured at different heights in the chromo­ sphere. The intensities were determined, from micro­ photometer tracings and characteristic curves. Corrections were made for the differential effects •

of the atmosphere, the instrument5 and the photo­ graphic plate response. The continuum was subtracted by assuming the photosphere continuum was due to a black body radiating at 5700° 1C and the extreme edge continuum resulted from the radiation of a black body at 4700°K. The numbers given for some of the Intensities result from overlapping the data from t' the two spectroscopes used. The data presented first are for a normal region on the sun. Data for an excited region are given later. Athay, Menzel, and Orrall (1957) found systematic variation from the intensities given here when they later remeasured CM’s plates. These differences will become apparent 131 in the comparisons of the published sets of {P } r given in the two publications.

AMO - 1936 eclipseo Athay* Menzel5 and Orrall (1957), recalibrated the data from Hemmendinger’s previously

mentioned 1939 dissertation in which he described height-dependent line intensities for 14 points - 9 of them normal, Menzel1s unpublished 1939 data for two additional points on the limb are also considered. Their data reduction gave results in a form consist:

tent with those from the 1952 data, ■7(19^1) - 1941 eclipse, Vyazanitsyn (1951) made use of a chromospheric arc technique to obtain his data, V(19^5) - 1945 eclipse, Vyazanitsyn’s (195®) technique was the same as for the 1941 eclipse, Athay and Thomas (1957) criticize the large spread of his data and find his results questionable, K - 1945 eclipse, Kristenson (1955) reduced Lindblad’s data (8 5 see, of film) to find good agreement with the linear log I vs, h curve at all but the lowest observed heights, Absolute intensities are not given, but [0 } are, For these low-dispersion data, K found that peak intensities were proportional to areas under curve, and so the peak intensities were used for the reduction. ABER - 1952 eclipse. Afchay, Billings, Evens, and Roberts (1954) produced the most extensive set of eclipse observation data to date. The data are reproduced and further analysed in Athay and Thomas (1957)» Areas under the profile curves were used to find intensities. An image divider with ratio IDOsl was used successfully in comparing lines of greatly different intensities. The effective height for each exposure was assumed to be the point reached by the moon’s limb at the midpoint of the exposure. The h == 0 point was estimated to be accurate to within ± 30km. The probable error given for the log of the intensities was ±0.04.

H(1952) = 1952 eclipse. Houtgast (1957) used narrow slit observation. Corrections made include consideration of the following wavelength-dependent features: atmospheric extinction, the transmission of the optics, and reflection coefficient. Because of timer-mechanism failure, he estimates that errors

are from 30$ - 100$ depending on the line considered. The sero point for h is uncertain, but results, such as they are according to H, are in good agreement with ABER. $$(1954) - 1954 eclipse. Houtgast’s techniques have improved. The slit width was varied for different. 133 sets of exposuresto add the advantages of the oa techniques to the ns ones„ The error in zero point is within 100km. Values given are averages for many observed intensities. Several normal regions on the limb were used for measurements. The maximum scatter in logarithms of intensities are given, as well as the number of points for which this scatter was considered. In the logarithm, scatter was on the

order of 0 .5 0 to 1 .1 0 for 5 ®r 6 points. MA - 1954 eclipse. Martynov and Abluseva (1962) used a method very similar to H (1954) except that their slit width was fairly wide and was fixed. The authors included a detailed account of the correc­ tions made, and they graphed comparison of their results with those of other observers. Agreement is fairly good. The discrepancy in the second and third limb-contact data is not discussed in the article but should be noted. St - 1961 eclipse. Airplanes trying to disperse the

cloud cover did not help Stolov (1 9 6 2) significantly.

He published values of , n = 4, 5, 6 . The data reduction was relatively careful, but the data were not very good.

CM excited - 1932 eclipse. A second point on the limb was studied by Olllid and. Menzel. TA excited - 1952 eclipse„ Ithay and Thomas (1957) reduced more data from the 1952 film. These data are repeated in Thomas' atidvAthay (1 96 1). Discussion of the Bata and Some Comparisons In this section the data in the following tables will be examined. It will be of interest 1) to see how well the data from a given article agree with themselves5 in order to determine their internal consistency; 2) to see how well one set of data compares with the others, in order to determine their mutual consistency; 3) to see how well the data fit the linear log^l(h) curve with which many investigators have tried to fit them, in order to determine justification of. this description; 4) to compile a weighted average of some of the data from original publications, and to use them; and 5) to prepare an average decrement which - will approximate the data within the limits of error for the data themselves. Weighting of data is intrinsically a subjective process. The remarks preceding the data and the results of examining points 1, 2, and 3 above will guide our assigning a weight to each set of data. Then we will attempt to carry out procedures 4 and 5® Figures 18 and 19 have been mentioned earlier„ In them the logarithm of the intensity of HP(=H4) has been normalized to log 100 - 2 for each decrement in every 135 Table 53• Balmer Decrements Published for the Solar Chromosphere without Limb Height Determination

1905* 1925* 1905* 1925* 1905 1926 1926 1920* 1937 1936 S DS M(1926) M(1947) T(1937; 3 (110) 100 692 200 4 (100) 100 100 100 5 90 (90) 80 6 9 .3 32 6 60 (8 0 ) 70 13 30 65 60 8 24 80 50 4 9 .0 9 24 60 45 2 6 .3 6.3 10 24 60 4o 14.8 5.0 11 21 50 35 8 .9 4.0 12 18 50 35 7.8 3.2 13 18 40 27| 6 .9 2 .5 14 15 40 25, 5.4 2.2 15 12 30 22i 4.3 2.0 16 12 30 20 3.8 1.6 IT 9 30 17| 3.2 1.4 18 6 30 15, 3.2 1.3 19 9 30 12! 2.3 1.0 20 6 30 12# 1.7 1.2 21 6 30 10 1.7 1.0 22 6 22 9 1.2 0.6 23 1* 18 9 1.0 0.6 24 18 9 0.9 0.5 25 14 6 1.0 0.4 26 14 6 0.7 0.3 27 10 0.5 28 10 0.7 29 12 0.5 30 8 0.8 31 12 0.3 32 6 1I 0.3 33 2 l! 0.2 34 10 1 0.2 35 2 1 0.2 36 2 0.3 37 0 0.2 136 Table 5^° Balmer Decrements Published for the 1932 Eclipse by 011114 and Menzel

h n 670 1500 2300 3170 4ooo 3 645 794 417 4 100 100 100 100 100 5 26=3 2 7 .6 21.4 25.1 39-8 6 13.2 14.1 10.0 11-7 O 7 8.7 H H 19.1 8 7.8 7*9 5-8 9-0 9 4.8 6.3 4.3 10 2.8 3.6 4 .0 2.6 11 2.5 3.2 3-1 2.3 12 2.3 2.9 2 .7 1.7 13 2.1 2.6 2.3 14 1.7 2.1 2 .1 15 1.4 1.7 1.5 16 1.2 1.4 1.4 17 1.1 1.3 1.3 18 0.9 1.1 19 0.9 1 .0 20 0.7 0 .8 21 0.7 0.8 22 0.6 0.6 23 0.6 0.6 24 0.5 0.4 25 0.4 0.4 26 0.4 27 0.4 28 0.3 29 0.3 30 0.3 31 0 .3 137 Table 55• Balmer Decrements Published for the 1936 Eclipse by Athay, Menzel„ and Orrall

h n 290 1000 1720 2440 3140 3850 4550 5250 5950 3 250 250 250 4 100 100 100 100 100 100 100 100 100 5 46.8 58.9 36.3 21.4 30.9 2 3 .5 1 7 .8 6 22.9 23-5 20.9 13-2 8-3 15-5 19-9 1 7 -0

8 14.5 14.5 11-5 9-0 5-0 9 7-8 6 .2 3-8 2-3 1-7 10 5,5 2.7 1.4 1 .2 11 5-1 1 .8 0 .8 12 4.7 1-5 0 .8 13 4-3 1.4 0 .6 14 3-6 1 .2 15 2 .1 0-9 16 1 .6 0.7 17 1.5 18 1-2 19 1-2 20 0.9 21 1.2 22 1.0 23 0.7 24 0 .7 25 0.6 26 0.4 27 0.4 Table 56. Balmer Decrements Published for the 1941 Eclipse by Vyazanltsyn

Locations Observed on the Sim t t ™: r— — . jx n 3900 4850 5T0o ^ 550 15'0'0 2600 3b0d -T5750~ 6900 3 1350 871 182 182 245 1095 4 100 100 100 100 100 100 100 1100 100 100 5 37.2 4 2 .7 57-5 31-6 5 6 .2 74.2 55-0 -38.9 39-8 3 0 .2 6 3 0 .2 3 3 .9 42.7 42.7 5 0 .2 3 0 .2 118.2 2 5 .1 T 2 6 .9 24.0 2 6 .3 33-1 33-1 22.4 1 6 .2 1 6 .6 8 24.5 18.6 15-5 . 23-4 25-1 1 5 -8 1 3 .2 9 9-3 19-1 7-8 10.7 16.6 3 .3

10 4.3 1.1 6.0 : 9-6 12.3 1-5 11 6.8 6.8 1-7 12 5-1 6.3 1.3 13 . 5-0 6.3 1.0 14 4.0 3-6 0 .9 15 2.6 2.2 0.7 16 2.2 2.0 0.7 17 '4::. 1 .2 1.0 0.7 Table 5 6 (Con'c.). Balmer Decrements Published for the 1941 Eclipse by Vyazanitsyn

Locations Observed on the Sun ~ W~"T9t>0 " 2 9 ^ 0 Woo" ™53o3^T5oo""'T!Oo^- 2150’.

3 174 263 159 162 123 1480 934 4 100 100 100 LOO 100 100 100 100 100 100 100 5 5 6 .2 57*5 38,0 38.9 38.9 5 1 .3 7 0 .8 2 8 .2 37.2 3 9 .8 33.9 -rt CO 6 42.4 40.8 18.6 19.1 22.9 3 3 .9 2 6 .3 23.4 2 3 .4 13.5 7 33.1 29.5 10.7 12.6 13-8 2 5 .7 33.9 22.4 15.5 17.0 7.4 . 8 23.4 20.4 6 .3 2 5 .7 19.5 11.5 9 10.7 19.1 3*2 1 0 .7 7.8 10 9.6 11.5 2 .3 1 0 .0 5.8 11 6.8 4.7 1.6 7 .6 4.3 12 . 5.1 4.1 1.3 7 .2 3.2 13 5o0 3*8 1.3 6 .3 3.0 14 4.0 2.5 0.8 4.0 1.9 15 2.6 1.7 0.3 4.1 1.1 16 2.2 . 1.1 0.5 2.9 0 .9 IT 1.5 0.4 0.1 1.4 0.3 18 1.1 0.2 0.5 0.2

H w vo Table 56 (Cont.). Balnte^ DecrementSI jPtibllshed for the 1941 Eclipse by Vyazanitsyn

Locatims Observed on the Stm r IT ‘4150 J20CT

3 275 2950 562 537 708 795 X4l 645 4 100 100 100 100 100 100 100 100 100 100 5 47.8 55.0 49.0 2 8 .9 7 0 .8 22.9 3 9 .8 1 9 .5 2.2 6 34.7 33.1 28.2 2 0 .9 2 2 .9 7 .2 T 27.5 21.9 19-9 14.4 1 5 .1 4.0 8 23.4 15.1 1 6 .2 2 .3 9 21.9 12.0 10 13-5 9.1 11 7.6 3.0 12 6 .5 2.7 13 6.3 2.7 14 4.0 1.8 15 2 .6 1.4 16 1.8 0.6 IT 0.9 0.5 18 0.2

H O Table 5 & (Cont.). Balmer Decrements Published for the 1941 Eclipse by Vy&zanitsyn

Locations Observed on the Sun h vrr n M o O ■'™69§ir . 4600 3600 'TOTT~ .. . ™ ”T O r 3 2450 3890 4900 3160 . 302 324 224 182 4 100 100 100 100 100 100 100 100 5 2 9 .5 3 1 .4 2 5 .1 5 .0 24.5 34.7 5 0 .2 6 1 3 .5 20.4 12.6 2.0 8.7 10.0 3 1 .6 7 7.8 14.4 7.8 1.1 4.9 6.6 2 1 .9 8 1 0 .0 5.5 0.7 3.3 5*2 1 7 .0 9 3.6 0.6 2.6 1*3 7*4 10 2.5 0.3 1.1 0.6 4.4 11 1.0 0.3 0.7 0.5 2.3 12 0.2 0.4 0.5 2 .0 13 0.2 0.2 0.4 1.6 14 0.1 0.1 0 .2 1.2 15 0.1 0.1 0 .1 0.6 16 0.1 0.1 0 .1 0.5 17 0.0 0.1 0 .0 18 0.0 0 .0 itri Table 56 (Cont.). Balmer Decrements Published for the 1941 Eclipse by Vyazanltsyn

Locations Observed on the Sun h ™ " 4960 ... . 13550”'™ 24bo lotr "iwo I (O) I (670)

3 457 398 257 288 355 4 100 100 100 100 100 100 100 5 47.8 24.5 27.5 46.8 24.5 57.5 55.0 6 24.0 13.5 12.6 28.2 14.5 40.8 38.0 7 13.2 7.4 7-2 19*5 9.5 •27.5 25.7 8 4.6 5*2 8.1 6.0 22.9 19.9 9 1.7 2.1 1.8 11.0 9.8 10 1.0 1.0 0 .6 7.8 6.6 11 0.7 0.8 0 .5 5.2 4.5 12 0.5 0.7 5»5 4.6 13 0.3 0.6 5.3 4.4 14 0.3 3.6 2.8 15 0.2 2.8 2.2 16 0.3 2.3 1.7 17 0.1 ' 1.1 0.8

- 1 8 . 0.8 0.4 142 Table 57o Balmer Decrementb Published for the 19^5 Eclipse by Vyazanitsyn

h point #1 2nd contact -* point #2 3rd contact ^ point #3 n : 650 1750 4250 5650 1000 2150 3450 4750 6200 4550 2150 600 I (0)

3 323 427 427 219 288 281 257 302 .4 3 6 457 :"302 4 100 100 100 100 100 100 100 100 100 100 V 100 100 :L100

5 67*7 7 8 .8 15.9 19c 5 74.2 15.9 67-7 33-1 2 0 .9 4 4 .7 51-3 61.7 6 52.5 35-4 6.8 5 0 ,2 18.2 24.5 8 .3 2 5 .7 41.7 51-3 *7f 8 16.2 2 2 .9 8.9 I3 .2 14.4 56.2 13-2 20.4 9 11.5 8 .5 4.3 4.5 5-4 9-3 10.2 10 8.7 4.2 1.2 3.4 2 .9 7-6 11 6.8 2.3 1.9 1.2 3-6 12 2.8 0.7 0.6 0.6 1.3

H W Table 58. Balmer Decrements Published for the 1952 Eclipse by Houtgast

h point b © •• point d n 132 553 125 536 970 1557 4 100 100 , 100 100 100 100 100 100 100 100 5 9 3 .3 5 7 .5 5 3 .7 3 1 .6 57*5 57*5 39*8 37*2 2 5 .7 3 2 .4 6 2 3 .4 1 6 .6 33*9 33*9 1 7 . 0 12.0 11.2 1 2 .3 7 2 7 .5 16.6 1 6 .2 7.0 15*1 11*5 6.8 5*3 4*5 5 .9 8 3 1 .6 14.1 1 9 .9 1 0 .5 17*4 11*5 10.0 6.8 . 7*8 1 0 ,1 9 1 3 .8 5*1 6 .9 2.7 6,3 4.0 10 6 .9 4.4 4.0 2.3 3*2 11 6.6 2 .9 3*7 . 1.9 4,1

Table 59» Mlmer Decrements Published for the 1954 Eclipse by Houtgast nnh’t -1000 0 1000 2000 3060 4ooo 5000 . 6000 7000 661 IO8 O 1350 1000 632 468 100 100 100 100 100 100 100 100 133 115 9 1 .2 64.5 52*5 51*3 5 4 .9 5 4 .9 I 39*9 40.8 2 2 .9 22.4 12,0 1 0 .5 9*4 31*6 3 1 .6 20.4 16.6 ■ 3*6 • 6.8 9*3 8 3 2 .4 Table 60. Balmer Decrements Published for the 1952 Eclipse by Athay, Billings, Evens, and Roberts

h (En) n 530 640 ": 750": 850 ; 960 ' ■■ 1070 r 1180 1280" 1399' 3 138 147 126 4 100 100 100 100 100 100 100 100 100 5 95.5 8 5 .I 100 88.1 7 9 .5 89.1 7 9 .5 84.2 6 9 .8 6 67.6 6 3 .2 6 9 .8 53.7 3 3 .9 5 0 .2 49.0 Y 3 8 .9 8* 38.9 20.9 ' 3 1 .6 24.5 1 6 .2 28.2 15.1 21.4 18.6 9 21 o 4 13.5 17.4 15.5 9 .3 14.8 18.2 1 1 .5 9-8 10 16.6 10.5 13.5 12.0 7 .4 11.0 4.9 8 .3 7-6 11. 13.2 9.0 1 1 .2 8.1 6.2 9.0 3.9 6 .9 5-8 121 13.2 8.5 1 1 .0 7.8 5.6 - 8.1 3.5 . 6 .0 4.6 13t 12.6 7.4 1 0 .0 7.0 5.2 8.1 3.2 5 .0 4.2 14 10.7 6.8 9.6 5.8 4.5 6.2 2.6 3.5 3-4 15 7.9 4.9 6 .2 4.0 3.X 4.3 1.8 3.0 2.4 16 7.5 4.1 5.9 3.2 2 .9 3.4 1.5 2.8 2.0 IT 5.2 3.4 4 .9 3.0 2.3 3.0 1.4 2.1 1.7 18 4.7 3.0 3.9 2.4 1.8 2.4 1.2 1.7 1.4 19f 4.3 2.8 3.8 2.6 1.7 2.3 1.1 1.6 1.5 20 3.8 2.6 3.3 2.1 1.4 1.7 0.9 1.5 1.2 21 3.9 2.6 3.5 2.1 1.0 1.6 0.9 1.2 1.0 22 2.8 1.9 2.4 1.7 0.8 1.3. 0.7 1.2 0.9 23 2.5 1.7 2.1 1.3 0.7 1.0. 0.5 0.8 0.7 24 „ 2.3 1.5 1.8 1.2 0.6 0.9 0.4 0.6 0.5 25y 2.3 1.4 1.8 1.1 0.5 0.8 0.4 0.6 0.5 26 1.7 1.1 1.4 1.0 0.4 0.7 0.3 0.5 0.4 2? 1.5 1.1 1.2 0.9 0.5 0.6 0.3 0.4 0.4 2 8 t 1.8 1.3 1.2 1.0 0.4 0.6 0.2 0.4 0.3 29. 1.1 0 .7 1.0 0.7 0.3 0.5 0.2 0.3 0.2 SO*1 2.3 1.3 1.5 1.0 0.4 0.6 0.2 0.5 0.3 145 31 1.6 0 .7 0.7 0.2 0.4 0.1 0.2 0.2 ^strong blend %f!aik blend Table 60 (Cent.). -BeCLmer Decrements Published for the 1952 Eclipse by Athay# Billings, Evens, and Roberts

h (ti) n 1500 1610 1730 1930 2260 2520 3060 3870 5490 6300 151 171 • 232 269 219 20? 724 1550 1515 100 100 100 100 100 100 100 100 100 ' 100 7 1 .4 81.3 66.1 6 3 .2 57-5 42.7 6 3 .2 4 9 .0 2 8 .2 I 3 7 .1 58.8 45*7 3 3 .9 3 0 .2 21.9 2 8 .9 2 3 .4 51.9 b 22.4 24.0 17.8 1 7 .0 14.4 8.6 10.0 9 12.0 12.6 7.4 7 .6 5-0 3-5 4.4 10 " 8.1 9.1 5-5 3-8 2.5 II 6.8 7.1 4.2 2.6 2.2 1:1 12 f 3*4 3*1 2.8 1-7 3.3 1 1:1 1:? 2 = 6 3 .0 1.8 1.3 1:1 14 ■ 3.2 4.1 2.2 1 .9 1.7 1.1 15 ■ 3.5 3.1 1 1.7 0.9 .0.9 16 2.2 2.6 1 4 1.6 0.8 0.7 1.7 2.0 1 2 1.4 0.7 0.6 II 1.6 1.4 1 1 1.0 0.6 0.4 1.4 1.7 0.9 1.0 0.5 0.3 20 1.4 1.1 0.7 0.8 0.4 21 1.1 1 .0 0 .6 0.6 0.3 22 0.8 0 ,7 0 .5 0.3 23 0.6 0 .6 0 .5 0,3 0,3 24 0.5 0 .5 0.4 0 .3 0.2 25? 0.4 0.4 0.3 0.3 0.2 '26 0.4 0.4 0.3 0.3 0.1 27 0.3 0.3 0.2 0.2 0.1 281 0 .3 0.3 0.3 0.2 . 0.1 29 0.2 0.2 0.2 0.1 '■ 0 .1 30l 0.2 0.3 0,2 0.2 0.1 31 0.2 0.2 0.2 0.1 146 tweak blend 147 Table 61. Balmer Decrements Published for the 1954 Eclipse by Martynov and Abluseva

Location h " " point a point b - n ':f3~ 1710 2210 148o 4 .100 100 100 100 5 41.7 28.9 6 13.8 : - 11.5 "15.1 19.1 7 8 8.7 7.8 8.5 11.0 9 8.9 5.4 4.2 9.0 10 4.0 2.1 2.0 4.1 11 3.2 1.8 1.6 3.0 12 2.9 1.5 1.4 2.6 13 2.8 1.4 1.5 2.1 14 2.0 1.1 1.1 1.7 15 .1.3 0.7 0,1 1.2 16 1.2 0.6 0.8 17 0.9 0,6 data from 2nd contact'

L o c a t i o n ______Averages 2nd 3rd point c point d point; e Contact Contact h n 3290 2740 2160 2610 2050 2490 1940 1(0) 1(0) 4 100 100 100 100 100 100 100 100 100 5 25.1 31.8 42.7 40,8 39.9 53.7 44.7 18.6 6 12.9 11.7 11.7 14.5 14.5 17.8 15.5 14.5 10.7 • - I 5.9 9.6 8.3 ; ' 7-8 9.3 - 3.6 9 1.5 3.3 4.2 4.2 8.9 3.7 10 1.2 2.1 1.7 1.7 4.8 2.6 11 1.0 1.2 1.1 1.5 4.1 2.2 12 0.9 1.2 1.0 1.0 3.7 2 .2 13 0.9 0.8 0.8 3.5 3.6 14 0.7 0.7 2.8 1.3 15 0.8 1.7 1.3 16 0.5 1.7 1.1 17 0.4 0.9 0.6 „ data frorn^rd contact Table 62. Balmer Decrements Published for Excited Regions during Eclipse

CM (excited) TA (excited) 1932 1952 n h ' 900 1730 2570 400 510 620 810 1350 1620 2160 4590 3 129 102 117 145 174 138- 4 100 100 100 100 100 100 (100) (100) 100 100 100 5 2 5 .7 40.7 ■33.1 72.4 64, 61.7 6 9 .2 8 3 .2 6 9 .8 6 11.7 20.2 7 12.6 8 10.2 1 6 .2 18.2 17.0 19.9 9 3.9 10.2 10.5 8.9 6 .9 10 2.2 7.6 7.4 5.4 4.7 11 1.9 6.2 4.2 4.2 3.8 12 1.8 5.8 3.6 3.7 3.3 13 1.7 4.8 3.9 3.3 3.0 14 1.3 3.8 3.2 2.4 1.9 15 1.1 3.2 2.0 1.8 1.6 16 1.0 2.6 1.6 . 1.4 1.4 17 0.9 2.0 1.4 1.0 1.0 18 0.7 1.7 1.0 0.8 0 .9 19 0.7 1.9 0.9 0.7 0 .7 20 0.6 1.3 0.6 0.7 0 .8 21 0.6 1.1 0.7 0.5 0 .6 22 0.5 1.0 0.5 0.4 0 .5 23 0.4 1.0 0.5 0 .3 24 0.4 0.6 0.4 0 .3 25 0.3 0.6 0.4 0 .2 26 0.3 0.5 0.3 0 .2 27 0.3 0.4 0.2 0 .2 28 0.3 0.4 29 0.2 30 0.3 H Figure 18. Some Representative Logarithmic Decrements of the Chromosphere Seen in the 1941 Eclipse by Vyazanitsyn. All decrements have been normalized so that log 1(4) = 2.00. The Greek letters designate the particular point observed by Vyazanitsyn on the limb and the height of the observation in kilometers above the limb. a'= point III at 850 km 6 = point VII at 1500 km P = point V at 2350 km e = point VII at 3600 1cm Y = point VII at 400 km C = point VIII at 2460 km 149

3

2

1

1

0

OI o 1

2

3 4 5 6 7 8 9 10 I I 1 2 13 14 15 16 17 18 n Figure lA. Some Representative Logarithmic Decrements of the Chromosphere Seen In the 1941 Eclipse by Vyazanitsyn. Figure 19. Some Representative Chromospherie logarithmic Decrements Seen in the 1952 Eclipse by Athay* .Billings, Evens, and Roberts, All decrements have been normalized, at log I*(h) == 2.0.

— Symbol. ^ Limb o 530 km « 850 km 4r- II80 km

fl 1500 km

0 1930 km a 2520 km § . 3036 km

t 5490, 1cm \) 63QO km o 3060 km A

II !

O) o -1

1— 1— 1— I— I— I I— I— I I I I I I I I I l i * i i » i i i i i i i i 10 15 20 25 30 n

Figure 19. Some Representative Logarithmic Decrements of the Chromosphere Seen in the 1952 Eclipse by Athay, Billings, Evens, and Roberts. 150 ; ’publleationi ■ The - normalized; logapitlimle ■dee$‘emente have been plotted vs. Balmer line number for several heights above several points on the limb. In this way# these two articles, v(l9^1) and ABER, were compared for relative internal consistency. The data from V(l9^1) are seen to have far wider scatter than those for ABER and far more irregularities in the shape of the decremental curve. In addition, in 7(1941) greater uncertainties are apparent in determining what decrement is to be associated with a given height. Other articles could also be considered by plotting them, in this fashion, but other aspects of the data may be compared to greater advantage. Figures 20 and 21 (pp. 153 and 15%) are also plots Of adjusted logarithmic Intensity data vs. Balmer line number. In these instances, however, one decrement is presented from each article in which the relative intensities of the Balmer lines are given as a function of their height above the disk. The decrements shown are for the values of h closest to 1000 km. The actual heights are presented on the graphs. Several conclusions may be drawn from these graphs. AMO said that they re-reduced CM’s plates and obtained substantially higher sets of intensities than OM did. The new f(3 } that they computed were greater by factors of about 1.5. The new intensity data are not published In 152 AMO’s article, but our graphs of the original CM data for both the normal and excited regions readily show that these original intensities lie considerably lower than those of later observers. One must note that the data of CM are given at lower heights in the chromosphere than those of other observers: however, this cannot be the source of the above discrepancy, for it alters the decrement in the wrong direction. If one accepts In(h) = ln(0 )e"’^nh, it is obvious that the decrement becomes steeper with increasing height. Thus, at h * 1000 km the CM values should diverge even further from the others. Some other irregularities are worth mentioning. A calibration error could account for the inversion of ABER curves at heights given as h = 960 and h = 107 km. The data of MA result in an abnormally steep decrement. Of additional interest are the irregularities in the values of V(l94l) (at h = 1100 km) and V(1945). The third and most extensive set of figures. Figures 22-27, pp. 156-161, contain all of the original data points for representative n-values (n = 3, 4, 5 , 8, 15). The logarithm of the normalized intensity ordinate is in what­ ever units were used in the articles (usually erg cirT^sterad”1 coming from a one-centimeter wide semi-infinite strip along a disk radius and commencing at height h, the height of the lunar limb above the solar). The abscissa is h. It is thus Figure 20. Logarithmic Decrements for the Chromosphere Near a Height of 1000 km. Normal Region. All data have been normalized at log 14(1000 km) = 2.0.

Height Inves­ Above Symbol tigation Solar Limb

• CM 670 km A AMO 1000 km 1 V(41) 850 km

D V(4l) 1100 km X V(45) 850 km 0 ABER 960 km

4 ABER 1070 km

none H( 52) 1020 km A H(54) 1000 km + HA 730 km -1 -

*— 1 1 1 I I I I 1 I I I I 1 1 I I I I I I I I 1 I I I I i I 1 1 I 5 10 15 20 25 30 n ----- ► Figure 20. Logarithmic Decrements for the Chromosphere Near a Height of 1000 km. Normal Region. 153 Figure 21. Logarithmic Decrements for Excited Regions of the Chromosphere Near a Height of 1000 km.

Symbol Limb Investigation Height

O CM [excited] 900 km D ETA [excited] 810 km A ETA [excited] 1350 km 3

2

C

0)0

5 10 15 20 30

Figure 21. Logarithmic Decrements for Excited Regions of the Chromosphere Near a Height of 1000 km.

H in -Pr 155 possible to see how closely the data are represented by a straight line. Kristensen's relatively calibrated curves are included for those n for which they are given. We find much variety in data from different articles. Because 1(0) values have sometimes been given, we have found that straight line extrapolations from other data points fall well above the measured values. For publica­ tions in which neither data for h = 0 nor an extrapolation is made, it is problematical as to how to make such an extrapolation. We have made none. Lines have been drawn between the data points for decrements without extreme data scatter. The broken lines may be compared with straight lines to judge to what extent In(h) = In(0)e~Pnh . Some interesting observations may be made concerning the data exhibited in this form: Intensities in the excited region of TA fade much more slowly with h than in any of the normal regions. The CM excited data, however, display a more rapid fading. This might be due to the aforementioned error in data reduction indicated by the discrepancies shown by the newer results for CM’s normal region. The greatest deviations from linearity occur for the low n-valued lines and for the lowest heights. On certain of the plots more accurate representations of the data points Figure 22. Observed Intensities of H3 as a Function of Height above the Solar Limb.

Symbol Investigation

1 CM 4 AMO + V (41)

X V (45) O ABER • H (54) A TA Excited 16

15

x:14

CO

13 U) o

12

0 2000 4000 6000 8000 h ( km )

Figure 22. Observed Intensities of H3 as a Function of Height above the Solar Limb.

VI o\ Figure 2 3 . Observed Intensities of H4 as a Function of Height above the Solar Limb.

Symbol Investigation

1 CM

1 AMO + V (41) X V (45) 0 ABER D H (52) # H (54)

4 MA A CM Excited A TA Excited K [Note: Kristensen gives no absolute calibration for his data.] 11— 1______I______I______I______1______I______L 0 2000 4000 6000 8000 h (km) ------►

2 3 . Observed Intensities of H4 as a Function of Height Above the Solar Limb. Figure 24. Observed Intensities of H5 as a Function of Height above the Solar Limb.

Symbol Investigation

1 CM A AMO + v (41) X V (45) o ABER D H (52) • H (54) A CM Excited A TA Excited K [Note: Kristensen gives no absolute calibration for his data.] i6r

013

0 2000 4000 6000 8000

Figure 24. Observed Intensities of H5 as a Function of Height above the Solar Limb

in oo Figure 25. Observed Intensities of H6 as a Function of Height above the Solar Limb.

Symbol Investigation

1 CM A AMO + v (41) X V (45) 0 ABER

0 H (52) # H (54) + MA A CM Excited K [Note: Kristensen gives no absolute calibration for his data.] 16r—

I 15

*>

SI 14 o

Q 1 3 * ♦

12

11 1 _J 2000 4000 6000 8000 h(km) -

Figure 25. Observed Intensities of H6 as a Function of Height above the Solar Limb. 159 Figure 26 Observed Intensities of H8 as a Function of Height above the Solar Limb.

Symbol Investigation

1 CM A AMO + V (4l) (2nd Contact) V (41) (3rd Contact)

X V (45) o ABER D H (52). • H (54) ♦ MA (1st Contact) MA (?nd Contact) A CM Excited

A TA Excited K [Note: Kristensen gives no absolute calibration for his data.] 16

oo

——K

12

11 0 2000 4000 6000 h (km)

Figure 26. Observed Intensities of H8 as a Function of Height above the Solar Limb. Figure 27. Observed Intensities of H 15 as a Function of Height above the Solar Limb.

K • CM A AMO + V (41) o ABER

% MA [2nd contact]

X CM excited D TA excited 161

1 6

1 4

1 3 cn o

1 2

j 1000 2 00 0 30 0 0 4000 height ( k m )

Figure 27. Observed Intensities of H 15 as a Function of Height above the Solar Limb. 162 can be achieved by a once-bent straight line than by an un­ bent one. The slopes, i.e., 6n ’s, above and below the bend would be different. It may have been noticed that among Figures 22-27, n = 7 was not included. Intensities are rarely published for this line, and similarly other lines seem also to have anomalous intensities. These lines are blended with neigh­ boring non-hydrogen lines just as the H line of Ca II dominates the H7 line. The resolution obtained by each author determined which of the Balmer lines could be iso­ lated. We indicate with the data of ABER, Table 60, p. 145, which lines they were unable to resolve. Using an exponential curve decaying with n to repre­ sent the intensity of a blended line, can present diffi­ culties if the two blended lines fade out at greatly different rates, i.e., if their (3's differ greatly. In this case it would be well to represent the intensity as the linear combination of two such expressions: 1(h) = However, since there was no indica­ tion that accuracies of intensity measurements are sufficient to warrant attempting a separation of blended lines, we have not considered this problem in our calculation of sets of average intensities. 163 The Average Chromospheric Balmer Decrement

An average Balmer decrement was calculated in the following manner: A) Each published source of data was assigned a weight by virtue of its favorable and unfavorable points (discussed in the previous sections). B) Weighted averages of the were computed, and height ranges were designated for which they are approximate­ ly valid. C) Weighted values for the logarithm of the intensity of each Balmer line at three different heights were computed. These heights are h * 1000 km, 2000 km, 4000 km. At these h-values the graphs of logl0In(h) were seen to be relatively linear. No computation has been made involving measured values of 1(0). This was avoided in the following manner. In general 1(h) = 1(0)e , and in particular I(h0 ) * I(0)e , for h = h0 .

Thus

or 1(h) = I(h0 )e'P(h*ho)

and ” 10Gio%(ho) — 2™T3 ^ ^ “ho ) D) Using the above calculation, three sets of values for log10In(h) were computed for fh}, where fh) = (500k kilometers} (k = 1, 2, . . . ) taken within the height range for which the line was actually observed. 164 By programming a computer to compute log10In(h) from the above equation, it was necessary only to feed in the values for { logl 0In( 1000km), f3n( 1000km) ) , { logloIn(2000km), 3n(2000km) ) , and { log,0In(4000km), Pn(4000km) } to generate the three sets of values. (The bars indicate average values.) E) These three sets of data were averaged and examined. P) From the same input information the computer was made to give relative Balmer decrements with I*(h) = 100. The manner of accomplishing this is simply derived. From above log,0l(h) = log,„l(h0 ) - (h-hy ) .

Now

= antilog {log In(h) - log I4(h)}

= antHog {log In(ho) - (h-h0) - {log I.(h0)

= antilog The values plugged In were, of course, the sets of averaged values used in the preceding calculation. To give 100, the final form was

Y ^ -y = antilog {log In(h0) - log I4(h0)

fh0} - {1000km, 2000km, 4000km) .

G) These three sets of values were averaged by the computer. H) The data were plotted and examined. The values, the plots, and interpretive remarks follow (l) the table of weights assigned the data, (2) the values given in the literature for {Pn), (3) the average f^n(h0j}, (4) the sets of (ln(h0)}, (5) a table of average logarithms, (6) graphs containing some of the computed logarithmic intensities, some of the averages of these intensities, and some average values computed and published by J . Houtgast (1962), (7) three tables of average Balmer line intensities and the average of the averages, and (8) some graphs of the non-logarithmic decrement plotted at four different heights, including a comparison of the 3 weighted average values of the decrement with the computed average decrement. 166 The following additional information is included in Appendix D: The Fortran programs for the above calculations, and complete tables of average logarithms that were computed. It was decided to weight the articles as shown in Table 6 3.

Table 6 3. Weights Assigned to Publications in Order to Compute (3n(h0) and In (h0 ) The number in parentheses indicates the number of h0 1s (if more than one) for which sets of data are given.

Information used Source Weight Pn (h0 ) In (h0 ) AHER 4 <*> |3| AMO(1936) 3 K 3 AMO(1932) 2 11(1954) 2 (3) MA 1+1 ♦ (2) 7(1941) 1 3) 7(1945) 1 (3) St 1 CM 0 / H(52) 0

♦2nd and 3rd contacts.

The average weighted intensities are given in Table 64 with the total weight in parentheses. The heights between which these lines are found in the chromosphere follow in the last column. The average of three computed sets of intensities is given in Table 66. The three sets of data are included in Table 64, Average Weighted Logarithmic Decrements for Three Representative Limb Heights

n log! In ( km) l°gioIn(1000km) 0 2000 logi©ln (4000km) h.min- h.max

15.58 15.37 (11' 14.71 (11] 1500km, 8000 km I 15.19 14.86 (12 13.85 (11 5 0 0, 7000 1 5 .0 8 14.57 (1 2 ' 13.58 (.11 500, 6000 I 14,82 I 14.35 (1 2 ' 13.25 (ill 500, 5500

8 14,57 11 ' 14.09 13.09 500, 4ooo 9 14.31 II 13.74 12.77 50 0, 4000 10 13.90 11/ 13.30 500, 3500 II 14.04 '11 13.32 $11 500, 3500 12 13.95 11 1 3 .2 2 11 500, 3500 14.03 ’10] 13.31 10] 500, 3500 11 13.90 '10 13.18 ;i° 50 0, 2500 13.66 '10 13.14 i° 50 0, 2500 % 13.68 10 ] 12.93 101 500, 2500 13.63 1 2 .9 8 6) 500, 2500 11 13.65 12.97 500, 2500 19 13.54 13.15 :1 ; 500, 2500 20 1 3 .6 0 Jl 12.99 4 500, 2400 21 13.41 ,7 1 2 .9 0 ‘4 50 0, 2400

22 13.32 , 7 11.86 ’4 50 0, 2400 13.23 ,7 12.78 '4 500, 2400 12 13.16 ,7 12.69 '4 50 0, 2400 25 13.18 , 7 12.62 ‘4 50 0, 2400

26 13.03 , 7 12.59 4 50 0, 2400 '4 27 13.13 J 1 2 .5 2 5 0 0, 2400 28 13.34 4 12.43 '41 50 0, 2400 29 13.18 '4 12.39 '4 500, 2400 30 13.32 '4 12.45 '4] 5 0 0, 2400 31 13.07 '4 12.26 ’4 500, : 2400 H

AKER AMO OM AHO TV V k n h==2000 • h= h ^ . (1936) (.1 9 3 2) (1949) (1945) 3 0.35 1 .3 0 1.20 1.16 0.90 0.86 . 4 0.87 1.60 •1.37 1.16 1.20 0.95 1 .3 0 5 =1.20 1.80 1.41 1.16 1.30 1.16 1 .3 9 6 1.36 . 1 .8 0 1.17 1.16 1.32 1.22 1 .4 3 T 1.18 1 .3 0 8 1.39 1 .3 8 1.20 - 1.42 1.01 1 .2 6 9 1.51 1,52 1.22 1.37 1.12 1 .5 4 10 1.56 1.39 1.22 1.45 1.12 1 .6 0 11 1 .6 2 1.82 1.22 1.46 1.15 1.66 12 1.68 1.82 1.22 1.47 1.11 1.70 13 1.68 1.87 1.25 1.52 1.91 14 1.70 1.88 1.25 1.8 1.60 I .85 15 1.73 1 .8 2 1 .3 0 1,6 1.55 . 1.82 16 1.73 1.93 1.36 1.9 1.68 1.22 .17 1.79 ' 2.0 1.28 2.0 1.70 18 I .85 2 0 0 1.20 1.9 : 2.04 19 1.91 '2,2 1 .3 0 • 1.9 20 1.91 2,0 1.39 2 : 0 21 1.91 2.4 1.42 2.1 22 1.96 2.4 1.33 2.2 23 1.96 2,2 1.47 2.3 24 2 .0 5 2.4 1.47 2.3 25 2.08 2.1 1 .5 8 2.7 26 2.05 2.0 1.64 2.2 27 2 .0 8 1.9 1.61 2.3 28 2.12 1.50 29 2.08 1.52 30 2.12 1.54 31 2.12 1.54 169 Table 65 (Pont.)„ Average Weighted Decay Coefficients hd) for Three Representative Limb Heights n

MA TA Weighted Average n 1st Contact 2nd Contact Excited [Normal Regions]

3 0.45 0.75 1 .0 6 (9) 4 1.45 2 .1 1 0.77 0.45 1 .2 0 1.40 (15' 5 1.75 1 .1 0 0.46 1 .3 2 1.4-9 (14] 6 1.34 1.95 1 .1 2 1.34 1.46 (15 7 8 1.45 1 .8 0 ■ I .3 6 (14] 9 1.54 2.04 0 .8 1 .5 2 (14 10 1.64 2 .3 4 0 .8 1 .5 5 (14 11 1.75 2 .3 6 0 .8 1 .6 9 (14 12 1,76 2 .3 0 0 .8 1.71 (14 13 1.84 2 .7 2 0.8 1,88 (13 14 ' 1.81 2.25 0 .9 1 .8 8 • (15 15 1 .7 6 2 .6 0 0 .9 1 .8 0 (15 16 1 .7 8 2.45 0.9 1.74 (15] 17 1.51 2.22 0.9 1.88 (12 18 0.9 1.92 (10] 19 0.9 2.0 (9] 20 1.0 2.0 (9 21 1.0 2.1 (9 ] 22 1.0 2.2 ' (9 . 23 1.3 2.1 (9. 24 1.2 2.2 (9 25 1-3 2.2 (9. 26. 1.2 2.1 , (9. 27 1.3 2.1 (9 28 2.1 (4 29 2.1 (4 30 2.1 (4 31 2.1 ' (4] Table 66. Average Logarithmic Balmer Decrements Computed for 14 Limb Heights

Height in Kilometers n 500 1000 1500 2000 2500 3000 3500

3 1 5 .6 0 15.42 1 5 .2 3 15.05 14.86 4 1 5 .6 9 1 5 .4 2 1 5 .1 4 14.87 14.59 14.31 14.04 5 15.55 15.25 14.95 14.65 14.35 14.05 13.75 6 15.27 14.97 14.67 14.37 14.07 13.77 13.47 8 1 5 .0 0 14.71 14.41 14.11 1 3 .8 2 1 3 .5 2 1 3 .2 3 9 1 4 .8 2 14.49 1 4 .1 6 1 3 .8 3 1 3 .5 0 13.17 12.84 10 1 4 .2 7 1 3 .9 4 1 3 .6 0 1 3 .2 6 12.93 12.59 1 2 .2 5 11 14.41 14.05 1 3 .6 8 13.31 12.95 12.58 1 2 .2 1 12 14.33 13.96 13.59 1 3 .2 1 12.84 12.47 1 2 .1 0 13 14.49 1 4 .0 8 1 3 .6 7 1 3 .2 6 1 2 .8 5 1 2 .4 4 12.04 14 14.36 13.95 13.54 13.13 12.72 15 14.18 13.79 13.40 1 3 .0 1 1 2 .6 2 16 14.06 1 3 .6 8 13.31 12.93 12.55 17 14.12 13.71 13.31 1 2 .9 0 12.49 » 18 14.14 13.73 13.31 12.89 12.48 19 14.21 1 3 .7 8 13-35 12.91 12.48 20 14.16 1 3 .7 3 1 3 .3 0 1 2 .8 6 12.43 21 14.07 1 3 .6 1 1 3 .1 6 1 2 .7 0 12.24 22 14.05 13.57 13.09 1 2 .6 1 12.13 23 13.92 13.46 1 3 .0 1 12.55 12.09 24 1 3 .8 8 13.40 12.93 12.45 11.97 25 1 3 .8 6 1 3 .3 8 1 2 .9 0 12.42 1 1 .9 4 26 1 3 .7 2 1 3 .2 7 1 2 .8 1 1 2 .3 5 1 1 .9 0 27 13.74 1 3 .2 8 12.83 12.37 11.91 28 1 3 .8 0 . 13.34 12.89 12.43 11.97 29 13.70 1 3 .2 4 1 2 .7 9 12.33 11.87 30 1 3 .8 0 13.34 12.89 1 2 .4 3 11.97 31 13.58 13.12 12.67 12.21 11.75 171 Table 66 (Conti.). Average Logarithmic Balmer Decrements Computed for 14 Limb Heights

Height in Kilometers n 4000 4500 5000 5500 6000 6500 7000

3 14.68 14.49 14.31 14.12 13.93 13.75 13.56 4 13.76 13.49 1 3 .2 1 12.94 1 2 .6 6 12.39 1 2 .1 1 5 13.45 13.15 1 2 .8 5 1 2 .5 6 1 2 .2 6 6 13.17 1 2 .8 7 12.57 12.27 8 12.93 9 12.51 172 the Appendix and were calculated from three sets of n pairs of constants: {(h)} = {(3,...,31),(3,•••,31),(3,•..,9)}• By examining the following figures, 28, 29, and 30, with some care it is possible to understand in detail what has been done in our computer calculations. For each n shown, three points have been indicated by a dotted line joining them. These points are from the tabulated set of log *n(h o (1 = 1,2,3); thus one point is the average intensity calculated for h0 = 1000, the second for ho = 2000, the third for h0 = 4000. A straight line is drawn through each of these points with its slope equal to the corresponding value of £n(h0). Each solid line thus represents the calculated intensity-values for log In(h) extrapolated from the intensity at one height h0 . The average logarithmic intensities from which the Balmer decre­ ments are calculated, are found by averaging the values of the three logarithmic intensities pointwisc at each height. These averages are plotted in Figure 30 for all observed n-values. The reliability of these averages may be estimated by noting in Figures 28 and 29 the variation within each set of solid lines, and also by noting how the dotted line (the actual intensities) relates to the solid lines (the calculated intensities). Except for the case of H3 the greatest divergence in the three curves extrapolated from hy-values occurs at the Figure 28. Computed Intensities of H3 and H4 as Functions of Height above the Solar Limb.

The .observed average intensities are indicated by the open triangles ( A ). The line connecting the two ^-8s gives the intensity varia­ tion with height computed from the average intensity at h = 1000 km. The line connecting the two x ’s gives the intensity varia­ tion with height computed from the average intensity at h = 2 000 k m . The line connecting the two o’s gives the intensity varia­ tion with height computed from the average Intensity at h = 4000 km. The open circles (o) indicate Houtgast*s (1962) average intensities. 16 16

I 15 I 15

14 14

CO 13 13 O) O) O o

12 12

0 3000 6000 0 60003000 h (km J -----► h (km) -----►

Figure 28. Computed Intensities of H3 and H4 as Functions of Height above the Solar Limb. 173 Figure 29* Computed Intensities of H5 and H6 as Functions of Height above the Solar Limb.

The observed average intensities are indicated by open triangles ( A ). The line connecting the two-l-'s gives the intensity varia­ tion with height computed from the average intensity at h = 1000 km. The line connecting the two x’s gives the intensity varia­ tion with height computed from the average intensity at h = 2000 km. The line connecting the two ofs gives the intensity varia­ tion with height computed from the average intensity at h = 4000 km. The open circles ( o) indicate Houtgast’s (1962) average intensities. I 5

14

-C

O 13 O' U) o o

12

0 3000 6000 0 3000 6000 h (km ) -----► h ( k m) ►

Figure 29. Computed Intensities of H5 and H6 as Functions of Height above the Solar Limb.

H Figure 30. Average of Computed Intensities for the Baliner Lines as Functions of Height above the Solar Limb. Three computed averaged for each Balmer Line were determined for heights above the chromosphere differing by 500 km. by methods explained in the text. The averages of these three averages are shown below.

Symbol n Symbol n

• 3 not skauin 18 + 4 net cWflun 19 X 5 rxet 20 A 6 not ckawm 21 V 8 not tVxOU* 22 1 9 not akown 23 10 24 T 11 not sKown 25 not sVxoUh 12 not 6Viowk 26 ■ 13 D 27 not s Wovn 14 not ckowun 28 *%#t aUoujk 15 not show* 29 + 16 not sUoun 30 A at lUo wn 17 1 31

For comparison, Houtgast’s average observed intensities as functions of limb height are given.

Symbol n

? 1 g ? 16

15

14

D) O 12

1 1 0 2000 4000 6000 8000 h ( km )

Figure 30. Average of Computed Intensities for the Balmer Lines as Functions of Height above the Solar Limb. 175 176 lowest heights (Figures 285 2 9 )» which is precisely what one expects after considering the slopes in the earlier Figures 22-27 (pp. 156-161). We observe in Figures 28 and 29 that each average of computed intensities at h = 2000 km must fall very near the actual value of the log In(h0) at that point. The computed average intensities will, in general, be somewhat greater than the actual log In(4000) and less than log In(l000) on the order of 0.2. At lower heights, i.e., down to 500 km, the computed averages become worse» At heights greater than h = 4000, the computed average is probably not as accurate an approximation as the one curve extrapolated from log !n(4000). This was the most linear portion of the curve in Figures 22-26. The other extra­ polated curves only reduce the average further from the actual values; however, as n increases from 3 to 6 in Figures 26 and 27 we see that the three log In(h0)-eurves are converging toward the upper limit of extrapolation. A more accurate averaging procedure would be to weight the data curves, so that the straight average of the three curves was used only between 1000-4000 km; thus we could have the appropriate log In(h0)-curve counting more heavily outside of that region. We do not, however, believe that the accuracy of our study would benefit substantially from such weighting. Moreover, the decrement would not be substantially changed. 177 The connected open data points in Figures 28 and 29 are Houtgast1s (1962) published averages as computed by averaging his own data with that from ABER, using other compilations as '’corrections." It is possible that his methods of weighting observers' results is more realistic than that accomplished here. It is not apparent, however, why Houtgast's decrements would be any more realistic; and so, for whatever use they may be, we have carried through these calculations to give a set of normalized average

Balmer decrements.

In Figures 31-34 we have plotted the average Balmer decrements for the chromosphere at four heights. In

Figures 32-34, we indicate the average of the observed inten- sity values. I.e., {ln(1000)}, Cln(2000)], and (iHptoOO)}.

It is possible to "feel" how consistent the averaging process was by noting how the correct intensity decrement value at each point compares with our average. We estimate that the variation from the average curve is generally less than 20^, and is far less for values of h within the region 1000 km s h £ 4000 km. Below 1000 km we recognize large uncertainty in our average— due primarily to lack of agreement of observed values and to departure from the simple exponential relationship. 178 Table 67. Balmer Decrements at Selected Limb Heights Extrapolated from in (1 dDd) and ^n ( n h 500 1000 1500 . 2000 2500 3000 3500 3 307.5 385.2 482.5 604.4 757.1 4 100.0 100.0 100.0 100.0 100.0 100.0 100.0 5 82.4 77.6 7 3 .1 68.8 64.8 61.0 5 7 .5 6 45.8 42.7 3 9 .8 37.1 34.6 3 2 .2 3 0 .0 I 26.0 24.0 22.1 20.4 1 8 .9 17.4 1 6 .1 9 15.5 13.2 11.2 9 .6 8.2 6.9 5 .9 10 6.1 . 5.1 4.3 3.6 3 .0 2.5 2.1 11 9.0 7.1 5-5 4.3 3 .4 2.7 2.1 12 7.4 5.8 4.5 3.5 2 .7 2.1 1.6 13 9.7 6.9 4.9 3.5 2 .5 1.8 1.3 14 7.2 5.1 3.6 2.6 1.8 15 4.0 3.0 2.2 1.6 1.2 16 4.0 3.1 2.4 1.8 1.4 17 3.9 2.8 2.0 1.4 1.0 18 4.1 2.9 2.0 1.4 1.0 19 3.3 2.2 1.5 1.0 0 .7 20 3.8 2 .6 1.7 1.2 0.8 21 2.6 1.7 1.1 0.7 0.4 22 2.2 1.3 0.8 0.5 0.3 23 1.7 1.1 0.7 0.4 0.3 24 1.5 0.9 0.6 0 .3 0.2 25 1.6 1.0 0.6 0.4 0.2 26 1.1 0.7 0.4 0 .3 0.2 27 1.4 0.9 0.6 0.4 0.2 28 2,2 1.4 0.9 0 .6 0.4 29 1.5 1.0 0.6 0 .4 0.3 30 2.1 1.3 0.9 0 .5 0.3 31 1.2 0.8 0.5 0 .3 0.2

' n h 4000 4500 5000 5500 6000 ‘ 6500 7000 3 948.3 1187.9 1488.0 1864.0 2334.9 2924.7 3 6 6 3 .6 4 100.0 100.0 100.0 100.0 100.0 100.0 100.0 5 54.1 5 1 .0 48.0 45.2 42.6 6 . 2 8 .0 26,1 24.4 22.7 8 14.8 9 5.0 179 Table 68. Balmer Decrements at Selected Limb Heights Extrapolated from In(2000) and ^n(2000)

n h 500 1000 1500 2000 2500 3000 3500

3 2 5 8 .3 323.6 405.3 507.8 6 3 6 .0 4 100.0 100.0 100.0 100.0 100.0 100.0 100.0 5 61.4 5 7 .8 5 4 .5 5 1 .3 48.3 45.5 42.8 6 38.1 3 5 .6 3 3 .1 3 0 .9 2 8 .8 26.9 2 5 .0 I 21.6 1 9 .9 1 8 .4 1 7 .0 15.7 14.5 1 3 .4 9 12.3 1 0 .5 8 .9 7.6 6.5 5.5 4 .7 10 4.7 3 .9 3 .3 2.8 2.3 1.9 1.6 11 6.0 4 .7 3 .7 2.9 2.3 1.8 1.4 12 4.9 3 .8 3 .0 2.3 1.8 1.4 1.1 13 7.8 5 .6 4 .0 2.8 2.0 1.4 1.0 14 5.8 4.1 2 .9 2.1 1.5 15 4.7 3 .5 2 .6 1.9 1.4 16 2.6 2.0 1.5 1.2 0 .9 17 3.7 2.6 1.9 1.3 0 .9 18 3.8 2.6 1 .8 1.3 0 .9 19 6 .5 4 .3 2 .9 1.9 1.3 20 4.5 3 .0 2 .0 1.3 0 .9 21 4.2 2 .7 1 .7 1.1 0 .7 22 4.5 2 .7 1 .6 1.0 0.6 23 3.2 2.0 1.3 0.8 0.5 24 3.0 1.8 l.l 0.7 0.4 25 2.6 1.6 0 .9 0.6 0 .3 26 2.1 1.3 0.8 0.5 0.3 27 1.8 1.1 0 .7 0.5 0.3 28 1.4 0 .9 0.6 0.4 0.2 29 1.3 0.8 0.5 0.3 0.2 30 1.5 1.0 0.6 0.4 0.2 31 1.0 0.6 : 0.4 0.3 0.2

n 4000 4500 5000 : : 5500 6000 6500 7000

3 796.7 998.0 1250.1. 1 5 6 6 .0 1961.6 2 4 5 7 .2 3 0 7 7 .9 4 100.0 1 0 0 .0 100.0 1 0 0 .0 100.0 100.0 100.0 5 40.3 38.0 35.8 3 3 .7 3 1 .7 6 23.3 21.8 20.3 18.9 7 8 12.3 9 4 .0 l8o Table 6 9. Balmer Decrements at Selected Limb Heights Extrapolated from 'ln (¥0'00')' and Pn (4000) n h 500 1000 1500 2000 2500 3000 3500

3 3 0 9 .3 366.7 434.8 5 1 5 .4 611.1 4 100.0 ■100.0 100.0 100.0 100.0 100.0 1 0 0 .0

5 73.6 7 0 .4 6 7 .3 64.3 6 1 .5 5808 5 6 .2 6 31.0 3 0 .1 29.2 2 8 .3 2 7 .5 2 6 .7 2 5 .9 7 8 15.1 1 5 .4 15.7 16.0 1 6 .4 1 6 .7 1 7 .0 9 12.7 1 1 .9 11.2 10.6 10.0 9 .4 8.8

n h 4000 4500 5000 5500 6000 6500 7000

3 724.4 858.8 1018.2 1207.1 1431.0 1696.5 2011.3 4 100.0 100.0 100.0 100.0 100.0 100.0 100.0

5 53.7 5 1 .3 49.1 46.9 44.8 6 25.1 • 24.4 23.7 2 3 .0 7 8 17.4 9 8.3 1B1 Table 70. The Average Balmer Decrement at Selected Heights These data were computed by averaging the corresponding data in Tables 67-6 9 .

500 1000 1500 2000 2500 3000 3500 5 ~ T 9iTf'“ SB". 5*" T 4 D . 9 a' W : T " ~ T 6r . T 4 1 0 0 .0 1 0 0 .0 1 0 0 .0 1 0 0 .0 1 0 0 .0 1 0 0 .0 1 0 0 .0 5 7 2 .5 68.6 64.9 ■6 1 .5 5 8 .2 5 5 .1 5 2 .2 6 3 8 .3 36.1 34.0 3 2 ,1 3 0 .3 2 8 .6 2 7 .0 1 2 0 .9 19.8 18.8 1 7 .8 1 7 .0 16.2 15.5 9 1 3 .5 11.9 10.5 9 .2 8.2 7.3 6 .5 10 5 .4 4.5 3.8 3 .2 2 .7 2.2 1 .9 11 7 .5 5.9 4.6 3 .6 2.8 2.2 1 .7 12 6.2 4.8 3.7 2 .9 2.2 1.7 1.3 13 8.8 6.2 4.4 3 .2 2.2 1.6 1.1 14 6 .5 4.6 3.3 2 .3 1.7 15 4.3 3.2 2,4 1 .8 1.3 16 3.3 2.6 1.9 1 .6 1.1 17 3.8 2.7 1.9 1.4 1.0 18 4.0 2.8 1.9 1.3 0.9 19 4 .9 3.3 2.2 1.5 1.0 20 4.2 2.8 1.9 1.3 0.8 21 3.4 2.2 1.4 0.9 0.6 22 3.4 2.0 1.2 0.7 0.5 23 2 .5 1.6 1.0 0.6 0.4 24 2.3 1.4 0.8 0.5 0.3 25 2.1 1.3 0.8 0.5 0.3 26 •1.6 1.0 0.6 0.4 0.3 27 1.6 1.0 0.6 0.4 0.3 28 1.8 1.2 0.7 0.5 0.3 29 1.4 0.9 0.6 0.4 0.2 30 1.8 1.2 0.7 0.5 0 .3 31 1.1 _ o / L _ 0.4 0.3 0.2

n h 4000 4500 5000 ■ 5500 6000. 6500 . 7000 3 8 2 3 .2 1014.9 1 2 5 2 .1 1545.7 1909.2 2359.5 2 9 1 7 .6 4 100.0 100.0 100.0 . 100.0 100.0 100.0 100.0 5 4 9 .4 46.8 44.3 41.9 39.7 6 2 5 .5 24.1 22.8 21.5 I 14.8 100

50

m

10 15 20 25 305 n

Figure 31• The Average Balmer Decrement in the Chromosphere at 500 km. A Balmer line intensities calculated for h = 500 km from average data at h = 1000. x Balmer line intensities calculated for h = 500 km from average data at h = 2000.

e Balmer line intensities calculated for h = 500 km from average data at h = 4000. o Average of Balmer line intensities calculated for h = 500 km. Figure 32. The Average Balmer Decrement in the Chromosphere at 1000 km.

x The average observed Balmer line intensities at h = 1000 km. A Balmer line intensities calculated for h = 1000 km from average observed data at h = 2000 km. • Balmer line intensities calculated for h = 1000 km from average observed data at h = 4000 km. o Average of the three Balmer line intensities at h = 1000 km given above. iue 2 TeAeaeBle ermn i h Chromosphere the in Decrement Balmer Average The 32. Figure 1000 km 100 50 0 at 1000 km. 1000 at 20 25 183 30 Figure 33. The Average Balmer Decrement in the Chromosphere at 2000 km.

x The average observed Balmer line intensities at h = 2000 km. # Balmer line intensities calculated for h = 2000 km from average observed data at h = 4000 km. A Balmer line intensities calculated for h = 2000 km from average observed data at h = 1000 km. o Average of the three Balmer line intensities at h = 2000 km given above. iue 3 TeAeae amrDceet nte Chromosphere the In Decrement Balmer Average The 33* Figure 2000 km 100 50 0 5 10 at 2000 km. 2000 at 15 20 530 25 184 Picure 34. The Average Balmer Decrement in the Chromosphere at 4-000 km.

x The average observed Balmer line intensities at h = 4000 km. • Balmer line intensities calculated for h » 4000 km from average observed data at h * 1000 km. A Balmer line intensities calculated for h = 4000 1cm from average observed data at h = 2000 km. 0 Average of the three Balmer line intensities at h = 4000 km given above. iue ^ Te vrg amrDceet nte Chromosphere the in Decrement Balmer Average The 3^• Figure 4000 km 100 50 0 5 10 at 4000 km. 4000 at n 15 20 25 185 30 186 Looal Phenomena

Definitions Bather than divide the material which follows into discrete sections for flares and for prominences, it was decided to put them in one section noting the differences of first quiescent and active prominences * then the similari­ ties between some very active„ brighta short-lived promi& nenees and certain active flares, and finally noting differences between limb and disk flares. Moustaches are mentioned in passing. These are small regions on the sun of lifetimes ranging from minutes to a few hours and associated with some bright continuous emission grains. Smith and Smith (1963) note that Balmer . lines to H10 have been observed. Ho measurements of the relative total emission intensities or intensity decrements have been found in the literature, so we shall not consider them further. Prominences originate in the chromosphere and project upwards to average heights of 30,000=40,000 kilome­ ters 1 extreme cases reach, on one hand, no higher than the spicules and, on the other, up to a million kilometers above the chromosphere» Hoyle’s distinctions between prominence types have been mentioned (p. 115), Menzel*s more recent classification (summarized in Smith and Smith, 1 9 6 3) is . 187 based on the motion of prominences and their association (or lack thereof) with sunspot groups. The quiescent 9 non-spot associated prominence has the greatest longevity, eight or nine solar revolutions, and undergoes relatively slow change. . On occasion, however, these types of prominences have been observed to change form very rapidly during an eruption that may send them to heights of up to 1& million kilometers above the limb. They then suddenly disappear, usually only to reappear in their original quiescent form within a few days. The active prominences are associated with centers of activity. Including spots, plages, and flocculi, on the sun’s surface. They take different forms, among which the surge and loop types are the most energetic. The duration of surges is on the order of minutes; of loops, hours. Published data useful to this study are neither so prolific in extent nor so unimpeachable in quality that extended discussion of individual variations in active prominence types is warranted. The data for all of the active promi­ nences will be grouped together, and when the description of one in any article is more explicit, this fact will be noted. In certain cases, it is a problem to know when a prominence is no longer to be considered a prominence, that is, to know that a flare is being observed. Smith and Smith point out that in the strict sense ”a prominence is a ' . 188 temperature and density anomaly In the corona which can be seen by its hydrogen emission on the limb or by its absorp­ tion on the disk,” whereas "a solar flare is primarily a transient localized brightening of the geometrically thin chromosphere.” They are "careful to distinguish between the flare and associated prominence-like activity. Even this distinction is not exact, for flares often acquire radial extensions into the corona, which are not always easy to differentiate from prominences/1 They point out that it is advisable to distinguish and separate the phenomena. In treating the published data we have tried to do that here to the extent that the distinction made by the researchers has allowed us. Terminology does cause some confusion. Hale, start­ ing with his first (1 8 92) description of flares, called them eruptions. Richardson (1939) objected because of the conno­ tation of ejecting matter and instead used bright chromo- apherlo disturbances. The word flare was apparently used first in English by MoNish (Richardson 1944) and is now generally used in English, whereas the French and Germans still use eruption. . Flare observations of the Balmer line intensities are more meager than similar observations for prominences. Richardson and Minkowski (Kuiper, ed. 1953) found the investigators of flare spectra to be "in a position somewhat 189 similar to a man trying to photograph the corona without knowing when an eclipse will o c c u r The fact that flares have been seen against the disk as well as on the limb is not yet an essentially useful facta for it is uncertain that the disk and limb flares are truly the same phenomenon. The problem is essentially a geometrical one in which effects at different layers of the solar atmosphere must be somehow separated in analysis into regions of physically different occurrences. A comparison between disk and limb flares is tabulated in Smith and Smith (1963, p. 1 2 8 )» Me mention some aspects of this tabulation concerning Balmer emission observations in our data section. Flares are classified by their corrected area on the sun’s surface (see* for example * Smith and Smith* p. 112). The largest flares are of importance class 3+ and the - smallest of importance 1-. Presentation of Data and Comments ' - The earliest recorded Balmer emission line inten­ sities found for prominences are those of K. Schwartzschlld (1906) and of Davidson and Stratton (1926). Schwartzschlld’s data for his observation of a "protuberance" came from the same eclipse plates that were reported in our chromosphere section. A similar procedure gave Davidson's and Stratton’s eye-estimated Intensities. The data are of low quality. It is unclear as to which type of object was observed. 190' Table 71• Decrements from Early Observations of Local Phenomena

Davidson1s and Schwartzschild’s mean n Stratton’s estimates values for 4 objects 4 100 5 35 6 40* 10 7 . 6 8 40 9 30 10 . . . 30 11 30 12 25 13 20 14 ' 20 15 ... . 15 16 13 17 12 18 10 19 9 . 20 . 8 21 7 22 6 23 6 24 5 25 5 26 4 27 4 28 3 29 2 30 2 31 2 32 1 33 1

*We arbitrarily set l(H6) = 40, 191 Prominences Ellison and Reid (1957) published equivalent widths and Balmer decrements for thirteen quiescent prominences which they studied with some care outside of eclipse. Their dispersion was high (better than 3 A/mm) and their observa­ tion times ranged from 5 seconds to 3i minutes. Corrections were made for grating ghosts, transparency, and darkening. One of the most interesting facets of their results is the linear relationship found to exist between central intensity and equivalent width for a given Hn line observed in different prominences; however, this relationship does not hold for H3- The resulting ratios found for central intensity (0.1.) divided by equivalent width W (X) are tabulated.

Table 72. Ratio of Central Intensities to Equivalent Widths in Quiescent Prominences as Reported by Ellison and Reid

H4 H6 H7 2.06 2.20 2.22

Zirin and Taudberg-Haussen (i9 6 0), in studying prominences, distinguished four classes of spectra. The characteristics of these do not, however, hinge upon the relative strengths of the Balmer lines which are relatively insensitive compared to the dramatic changes undergone by 192 weaker metallic lines. In examining a particular quiescent prominence, the observation was made that the spectrum is "exactly that of the chromosphere at 1500 km." The remark probably was meant to point out that the lines present in the two spectra are identical, for it is very apparent from comparing the appropriate column in Table 69 (p. 180) with Zirin ahd Tandberg-Hanssen in Table 72 (p. 191) that the Balmer decrements are quite different. In a paper on prominence spectra, Jeffries and Orrall (1962) obtained, high dispersion data during a period of "good seeing conditions." Intensities for two prominences were calibrated carefully In the limited spectral region observed, XX36OO-38OO, in hope that photographic errors could be avoided. Unfortunately direct conversion of the data to a form useful for comparison in this study Is not easily accomplished. If an overlap of these data and those of Ellison and Reid occurred, it would be possible to normalize the Jeffries and Orrall data to some intensity for the lines common to both studies. Since the two data sets have zero intersection, the only possible comparison would be that of the deeremental slopes in the two. regions. Me, however, found no linear relation when we plotted the inten­ sities as functions of n and of X. Jeffries and Orrall have chosen to plot intensity vs. log where A Is a con­ stant . • The slopes from averaging values of Intensities for 193 two prominences may result in a straight line— but the physical meaning of such a procedure is not apparent} and we have not found this procedure applied elsewhere in the literature for any other Balmer line intensities. Perhaps it is worthwhile. Regarding small differences between the two objects they viewed# Jeffries and Orrall associated them with "insuring that the microphotometer slit traced through exactly the same place in each spectral arc# father than to effects of self-absorption or departures from equilibrium in those high lying levels Jeffries and Orrall1s data appear in Appendix E. Also In Appendix E are the unnormalized data from publica­ tions in which the form of the decrement differed from that used throughout this thesis# i.e.# those with l(H4) = 100. Active Prominences Published Data A. UnsBld (1947) remeasured and re-analyzed the spec­ trum of an "eruptiven Protuberanz11 (hopefully an eruptive prominence and not a flare) that was taken during the 1929 eclipse. He measured and reduced one of six exposures# and explained in detail his attempted corrections for various factors. We find his work is crude by later standards. Most observations of active prominences have been recent. Table 73* Published Balraer Decrements for Quiescent Prominences

Ellison and Reid (1957) Prominence Number 1 2 3 4 5 6 ' 7 8

3 243 332 358 762 370 256 1111 • 607 4 100 100 : 100 100 100 100 : 100 1 0 0 . 5 33*2 29*4 1, 2 7 .6 4 5 .3 3 2 .3 39.4 2 9 .5 3 0 .3 6 7*5 0 .4 • X: 5*8 , 1 8 .7 Calculated from equivalent widths

: lEHMoh- :and;.Reid2.41957) .. Zirin arid ; . Number Taudb erg-Hdussen 9 10 11 12 13 Average (I960) 1122 726 , 5:408 535 563 5 1 1 100 100 : 100 100 100 100 ; 100 X. x 2 3 .5 45*7, 45-3 45*1 48.1 I 2 7 .3 X: 2 5 .0 33*1 n* 47*0 14.0 , 12.7 15*5 14.1 8.0 I 6.7 Calculated from equivalent widths Calculated from intensities 195 Athay and Orrall (1957) obtained speetra of a promi­ nence at the 1952 eclipse, ' They subsequently reduced the • data by Integrating the area under the line traced by a mierophotometer to get the total emission,

Zirln and Tandberg-Hanssen (i9 6 0) compared their ' quiescent promlneiiee data with observations of two active prominences which accompanied a huge flare of importance 3+ on the limb. They found that both regions had exactly the same spectra as "every loop, surge, or flare that we have ever observed (about 25 oases). It is quite remarkable,” they comment, "that the scatter in line intensity ratios is small among these objects, even though these ratios differ greatly from those in quiescent prominences." Two years prior to the preceding observation, Elliot, Ellison,'and Reid (i960) observed, a similar promi­ nence which accompanied a flare of importance 3- Their measurements were for both the equivalent width and central Intensity (measured with respect to the center of the disk).

Dispersion was about 3 X / m m In the first order. Since H4 was not measured we could not compute a normalized decrement3 however, we have presented the data here since they may readily be compared to the results of Zirln and Tandberg-

Hans sen by noting the H 6A l 7 ratio. We cannot explain the discrepancy. 196 Jeffries and Orrall (1961a* 1961b) Investigated the spectra of a loop prominence and a spray-type limb event with the same care that they exhibited In their observations of the quiescent prominence (of p . 179). As with their quiescent prominence observations the Balmer series for the active prominences begin with Hll. Because we are unable to compare these results meaningfully with the other data> they are displayed in the Appendix. We have already stated that both of these objects are of untypically high energy, and we shall at least be able to make a comparison' with data appear­ ing in Table 73» Yeh Shih-Huei (1961) published uncorrected inten­ sities from ten. prominenceso The data are seen to vary by factors of 4 or greater in more than half of the Balmer lines observed in the prominences. Their average must be weighted very lightly. Analysis of the Data It is interesting to note that the average intensity values obtained by XJnsbld, Athay and Orrall, and Shih-Huei agree well between H4 and H6, beyond which they diverge rapidly for the different observers. The average inten­ sities for these three early Balmer lines are given in Table 75. Beyond the height of 5000 km in the chromosphere the Balmer lines beyond H6 are not usually seen, but the Table 74. Published Balmer Decrements for Active Prominences

TTncjRi^ Athay and Zirin and : : "Elligtt^ Ellison,, and Reid Orrall Taudberg-Haussen W 1(X) I 122. 105 3 775 528 200 4 100 100 100 5 47.8 4 4 .7 6 22,4 2 1 .3 33 36 45 7 2 8 .9 6 39 52 8 25.7 1.8 1 9 8.5 1.1 10 3.5 1.0 11 , 3.2 1.1 12 3.6 0.8 13 4.1 0.6 14 4.0 0.4 15 1.2 0 .3 16 1.7 0.2

H IQ -4 Table 74 (Cont»). Published Balmer Decrements for Active Prominences

10 Prominences Observed by Shih-Huei 3 520 250 250 290 610 600 510 1510 620 530 570 4 100 100 100 100 . 100 100 100 100 100 100 100 5 31 39 46 27 46 61 54 27 55 56 46 6 16 36 12 18 30 24 46 14 17 30 24 7 11 25 T«5 13 24 18 . 43 • 11 13 22 19 8 6 14 4.8 4.8 11 . 7-4 15 2 .0 5.2 9-6 8 9 12 2.7 2.7 3-7 3-6 12 2 .6 5-4 6 10 2.7 3-1 13 1-9 3-9 5 11 3-4 3-0 8 .1 2 .0 3-7 4.0 12 3-1 4.5 5-1 1-3 2 .0 3-3 13 1-9 2 .1 6 .2 1-3 1.9 2.7 14 1 .0 1.5 1.3

H VO 00 199 Intensities for H4-H6 at h = 5000 km are displayed below for comparison.

Table 75• A Limited Average Balmer Decrement for Active Prominences

n 4 5 6 3-article average 100 46.2 22.6 average chromosphere at h = 5000 km 100 44.3 22.8

Although the Jeffries and Orrall measurements are for high n we were able to compare slopes of the decrements for their highly excited objects with the decrementa1 slope of Athay and Orrall*s limb event. Jeffries and Orrall took series of exposures covering relatively long segments of the prominences’ time on the limb. The decrements changed during this time, so that we were able to observe development with time of the slopes of the logarithmic decrement between Hll and Hl6. Athay and Orrall found log 1(11-16) = log I(Hll) - log I(HI6) = 0.63« Jeffries and Orrall found the values given in Table 76„

Table 7 6. Slopes of the Logarithmic Decrements in two Active Prominences

log 1(11-16) Average

spray-type event 0.22 0.28 0.53 0.31 0 .2 6 0.26 0.29 loop prominence 0.21 0.42 0.26 0.20 0.27 200 The change with respect to time is probably real. For both these events the decrements are considerably shal=- lower than the Athay and Orrall event. This is consistent with the assertion that the loop and spray-type events are very energetic events, and also with the observation that the Balmer lines are visible to H25 for the quiescent promi­ nence observed by Athay and Orrall. In general, the correlation among the sets of promi­ nence data does not seem high enough to take any general average or to justify our attempting to give typical data - for different types of events. Limb Flares The intrinsic problem with observing limb flares is observation of the flare Itself. The geometry must be right in order to record a flare phenomenon? that is, l) the limb must not obscure it by being between the event and the observer? 2) any part of the limb which lies behind the flare must be corrected for. For a truly unobscured limb flare one must determine the limb height at which the flare is occurring and concurrently must try to avoid observing associated eruptive prominences and the chromosphere. Due to the vast intrinsic observational differences, it is uncertain just how flares seen onthe limb relate to those viewed against the disk. 201 Richardson and Minkowski (1939) published inten­ sities of bright ehremospheric eruptions viewed on the limb from H3 to H8 . These were based on eye-estimates of rough photographic spectra. Allen (1940) published an average set

of eye-estimated intensities from H3 to WJ compiled from 300 . "eruptions <.’* The extra-eclipse observational mode used for these investigations added scattered light uncertainties t© the photographic calibration difficulties«. Richardson (1950) observed a more recent flare and published eye-estimated intensities. Neither these inten­ sities nor those published earlier contained any corrections • made t© the raw data» A more serious effort was; made by Polnpan (1960)3 who observed a limb flare photographically and reduced it by.;; microphotometer bracing to obtain equivalent widths.,. He then corrected for limb darkening;. in order to reduce these :t© intensities at the,-center of the disk. Bisk Flares ; ; - Flares, seen against the disk seem highly irregular ' with respect to their . Balmer decrements. . Smith and Smith (1961) comment?: As a general rule 5 when the E® central inten-, : sity is only slightly above that of the local con- fcimmm* the^ other Balmer lines, including He and Hg have slightly lower intensities than However, in bright flares,,, where H® approaches and exceeds twice the local continuum Intensity, the behavior of the central intensity of the later lines, and 202 Table 77» Published Balmer Decrements for Limb Flares

- :-v Riohardson Allen’ ■--- ’ and" Minkowski" " Riohardson Polupan 3 4oo 400 930 329 4 100 100 100 100 100

5 125 250 50 75 ' 47 •6 100 34 70 22 7 25 34 .50 14 8 26 45 5 9 40 10 35 11 30 12 32 13 25 14 25 15 20 16 8 17 ; 9 - 18 : 5 19 3 20 3 21 " 2 203 particularly HP to H8, may vary considerably from flare to flare and within the same flare. HP may "be brighter than Hct, while Hy and later lines are less intense than H

Due to poor viewing conditions ^vestka (i9 6 0) was able to take only one spectrum of a flare of importance 2. 204 The equivalent widths were found after subtracting background and correcting for scattered light. The corrections depend on using a source function based on an assumption of thermo­ dynamic equilibrium. A decrement was thus derived. The original data are given in an earlier article describing the apparatus (Valnl&ek, Latus, Blaka, Svestka, and Seidl 1959). Smith (1963) differs substantially with Svestka for another disk flare of importance 2. Despite observational difficulties, seventeen sets of relative intensities cover­ ing a period of 47 minutes were obtained. Clouds and coelostat guiding were problems and the photometry was complicated because the flare overlay a sunspot. The data given here are central intensities above the continuum level. Because of variable widths of the lines, these intensities can not give a true decrement, i.e., the absolute inten­ sities ratios remain uncertain. She suggests that a source function is implied by the data which increases with height of the flare above the disk, i.e., she infers a highly excited surface exists locally. Source functions are discussed along with other aspects of the observational problems in Smith and Smith. The usual flare equation is presented.

Ix = IBe'TX + S(l-e*TX), where is the background Intensity; is the total optical thickness of the flare in the line of sight; o’ is the mean Table 7 8 » Published Balmer Deerements for Flares Seen Against the Disk

Jeffries j, Smith,, and Smith (1959) Pre~ Post- Pre- Post- Valnieek ^ Maximum Flash Maximum Maximum Maximum Flash Maxiimim Maximum et al Svestka 3 125 115 78 79 330 120 175 64 56 89.4 182 4 100 100 100 100 100 : 100 100 100 100 100 100 5 90 125 78 90 100 75 60 55 48 71.3 9 2 .6 6 90 75 54 66 83 . 67 35 30 44 7 75 69 43 61 54 33 "S' 73 0 5

Smith (1963) for 17 flares 3 115 1:115 I? 138 :135 117 115 64.6 52.5 4 100 100 100 100 100 100 - 100 : ; 100 100 5 95.5 91.2 8 9 ,1 91*2 9 1 .2 7 0 .8 58.5- 5 2 .5 5 0 .1 6 8 7 .2 8 5 .I 9 3 .3 8 5 ,1 77.6 79 0 4 41.7 46.8 5 8 .9 7 > 68.1 83.1 8 1 .3 8 5 .1 ',- : 74.1 .. 8 3 .1 25.1 39.8 5 7 .5 8 7 6 .0 79 »4 66.1 . 8 1 .3 64.6 64.6 , 67.6 45.7 5 1 .3 ‘Sable 78 {Cent.), Published Balmer Decrements for Flares Seem Against the Disk

Smith (1963) for 17 flares Average

3 69 »2 251 141 117 162 159 8 5 .I 135 120 4 100 100 100 100 100 100 100 100 100 5 93.4 69-2 81-3 8 9 .1 7 7 .6 7 4 .1 44.7 77-6 76 6 74.1 72-5 63,1 87-1 5 8 -9 69-2 43-7 6 7 -6 69 7 77-6 46.8 72-5 81.3 64.6 6 7 .6 38.0 58-9 63 8 63-1 39-8 44.7 77.6 52.5 72-5 31-6 52-5 61 206 20? source function. In presenting source functions large disagreements exist in reports by different investigators.

It would be well to summarize the solar observational material before we conclude this chapter. This is, howevers not easily done. Although chromospheric data may be greatly improved, especially close to the limb, they were sufficient­ ly abundant to allow us to compute an average decrement. The variations of observed intensity values from the average were not excessive for heights above 1000 km beyond the disk. For numerical values we refer back to Table 70. The observations of the Balmer decrements in local disturbances, flares and prominences, might have been com­ pared with the chromospheric case in a fairly detailed and interesting fashion if we did not find such extensive inconsistencies in them. Generalizations may still be made if we note, for example, that the decrement of a disk flare is much shallower than that of the chromosphere; or that those of active prominences might be somewhat steeper. Even these generalizations are somewhat dangerous, inasmuch as measurements and computations for obtaining decrements in different types of solar phenomena are not necessarily equivalent. The problem of calculating from central Intensities Instead of total emission was described 2 0 8 earlier. In many cases we can not consider it as an accurate first-order approximation. Spreads in published intensities for local solar phenomena are indicated in Table 7 8 . As can readily be seen, better data are needed badly before definitive decrements may be meaningfully promulgated. Table 79« Baln©r Emission Line Intensity Spreads Published for Looal Solar Phenomena The number of objects included in each group is indicated in parentheses.

n Quiescent Active Limb Flares Disk" Flares Prominences (l4) Prominences (10) , (4) "7 243-1122 105-1510 2 9.5-330 4 100 100 1 0 0 100 5 27.6-48.1 27-61 47-250 48-125 6 0.4-47.0 12-46 22-100 30-90 7 8 -15.5 6 .6-52 14-50 2 5 .1-8 5 .1 8 6 .7 1 -25.7 5 -4 5 31.6-81.3 9 1 .1-8.5 4o 10 1 .0 -1 3 : 35 11 1 .1- 8.1 30 12 0 .8-5 .1 ■ 32 ' 13 0.6-6,2 ' ■ 25 14 0 .4-4.0 25

15 0.3-1o2 •- 20 16 0.2-1 .7 8 : 17 0 , ""-l. 7 ■ 9 18 5 . 19 ' 3 20 . 3 2 21 605 CONCLUDING REMARKS

Other astronomical sources of Balmer emission lines exist in addition to those for which decrements have been quotedo Some examples are notable. The Mira variable stars are often recognized by their outstanding Balmer line emission characteristics, but for them we could find no data published after our 1950 cut­ off date.

Galaxies

In 1943 Seyfert discovered that Balmer line emission was characteristic of a certain unusual group of galaxies which were subsequently named after him. Humason had given eye-estimations of intensities for one of these objects before Seyfert in 1932. We present the data from both studies in Table 80, while noting that we have found no subsequent Balmer line intensity measurements for members of this group of galaxies. NGC 7027 is not a spiral but a gaseous nebula; the others are spirals. Although Seyfert did make some observa­ tional corrections, his intensities would not be considered accurate today. Nevertheless, his work serves to point out this additional area where Balmer line intensities may be measured and decrements calculated.

210 2 1 1 Table 80. Published Balmer Decrements for Seyfert Galaxies

Humason (1932) Seyfert NGC 1275 NGC NGC NGC 1068 1275_.. 3516 4051 Ha + [Ne II] 1000 700 600 600 HP 100 ■100 100 100 100 Hy 313 40 50 4o 40 H6 63 20 10 25 He 12 25 20 HC + [He III] 5

Seyfert (1943) NGC 4151 NGC 7469 Coffe and ■ Core and . NGC Core: wing. Wing: / Core ' Wing 7027 Ha + [l;e II] 400 375 400 420 H P . 100 100 100 : 100 100 100 . H y 35 30 35 60 20 H8 20 . 20 20 35 12 He 25 5 15 HC + [Ne III] 7 Qmasl-Stellar Objects

Probably not all Balmer emission sources have yet been discovered. Recently * Gke (19'65) presented in his study of 30273s, the first published measurements of a Balmer emission decrement in a Quasi-Stellar Object. The measure­ ments were photoelectric and their resolution lows e.g. H6 already is blended with He.

Table Bl, A Published Balmer Decrement for a QSO

w f(x) Balmer Decrement

m 330 185 w 98 100 HS 29 38

Taking account of all his measuremental uncertain­ ties leads Oke to conclude that at worst l(H®} = 200. 185 Is his lower limit.

Summary

Thus it is seen that there are new objects to meas­ ure and old ones to remeasure. Data are being accumulated and are gradually improving* through the efforts of the astrophysicists who are attempting to take meaningful observations and calculations. The data that exist have 213 been given along with evaluation and. criticism where we felt It was appropriate a Present and future data needs were stressed« Average decrements for similar observations, in the case of the solar chromosphere, and for similar objects, in the case of the gaseous nebulae, have been computed, for their limited worth. fhe suggestion of some investigators has been repeated in this study that the highest accuracy data thus far available are the result of combining photographic and photoelectric measurements. Me think it is appropriate to conclude with another suggestion regarding techniques for improving measurements for rapidly occurring events, in particular, for chromosphere measurement during times of eclipse.

A Suggestion

Iv f.. "'.v.ltiscsuggested that .-a speotroscope':be ’constructed'; • utilising many tiny photosensitive diodes instead of either a photographic plate,or scanning phototube. During the course of an eclipse the outputs of these diodes could be recorded in an electronic memory device, which in turn could later be fed into a computer for analysis. Many advantages can be cited for such a device, ranging from minimising human operation and error to enabling an easy analysis of the data, from having fairly high line resolution without eeneurrenti photographic problems to having continuous instead of discrete variation In the height parameter» Considerable effort would be needed to make such a system operatives but its contributions might become invaluable. APPENDIX A

Balmer Bine Intensity Data as Originally Published for Nebulae

215 . 2 - , 1 -i 6 " Table 82. Balmer Line Intensities for Nebulae .■ Published,In Kal^y (1964) ,. -v . ,, I0 ~ observed Intensity! Ic '= Intensity after making a oorreotion for Intersteliar r e d d e n i n g . ____ Anon 21^31m 10 2165 NGO 3242 n :• IT X X_ 1 *0 "0 © . %e 3 4 100 100 100 V 100 100 100 5 46 49 . 46 : 46 63 63 6 22.5 24.6 26.5 1 26.5 32 32 7 17 19.1 ' 8 7*3 ■ 8.1; 18 k 13 - 9 5.35 6.2 • 6.3 :; : 6.3 7*4 7 .4 10 4.18 . 4.8 4.45 ;;ii'v 4. % 6 .5 6 .5 11 3 .0 6 3.5 3.7 ; 3.7 5.4 5 .4 12 2.58 2.95 2.9 :■ 2 .9 4.5 4 .5 13 4.42 2 .4 5 !: 2.45 3*5 3.5 14 1 .8 5 2.15 3 .2 r 3 .2 3.2 3.2 15 2.16 2.45 1 .8 ■ 1.8 2.3 2.3 16 1.33 I .52 1 .8 5 1 .8 5 17 1.20 1.38 1.4 i . 1.4 1.6 1.6 18 O.gO 1.03 1 .3 5 r ■ 1*35 1.48 . 1.48 : 19 0 .8 0 0.91 1 .1 ■ 1.1 1.32 .. 1.32 20 O.BO 0.91 1 .0 l-i':. 1,0 1.2 ' 1.2 : . 21 0.73 0 .8 3 0 .8 0 ?i . 0.80 0.91 0.91 • : 22 O .7 3 0.83 0 .8 0 : 0.80 0 .8 1 0.81 23 0 .7 0 0 .7 0 0.78 0.78 24 0.81 0.81 .

Anon I8h15m Anon l8h47^ Anon 22h29m r.-* j . . Io I© t • l e Io Io

100 100 loo : 100 100 100 5 51 51 46 46.8 52 . 52 6 26 26 . 25 25*7 35 35 7 8 11 11,5 25 25 9 8.1 8.1 7.7 8.1 10 10 10 5.0 5*0 5*4 7-1 9.0 9.0 11 4.7 4.7 4.1 4 ,3 6.2 6.2 12 4.7 4.7 3.7 3*9 5*4 5*4 13 5.5 5*5 14 2.6 2.6 2.5 2.6 3*9 3*9 15 2.2 2.2 1.9 2.0 Table 82. (Continued)

Anon 23E 24 10 4846 10 5 U T 10 5217 J 900 NOG 40 n L 1, I, o © :o I. o Q 3 "1^5— 319 4 100 100 100 100 100 100 100 100 100 100 100 100 5 36 49 49 49 42 49.7 55 55 42 48 43 48 6 17 ■ 26 27 23 26.3 28 28 21 25 24 28 7 11 42 8 11 • 19 21 21 12 14.5 20 20 14.8 18 9 6.0 11 8.8 8.8 5.1 4.9 6.0 6.0 7.4 9.6 11 1 3 .5 10 4.9 8 .9 IS 7.5 3.8 4 .7 5.2 5.2 5 .0 6.6 8 10 11 4.1 7.4 4.1 4.1 3.5 4.3 3.2 3.2 4.2 5.5 6 7 .6 12 4.0 7.4 4.9 4.9 2.2 2 .7 2.9 2.9 5.5 7.3 5 6 .3 13 4.8 9.1

MGC 4 3 OI MGO 6309 NGC 6720 HGO 6790 MOO 6803 MOO 6807 ti I, [„ I„ © "d 391 302 4 100 100 100 100 100 100 100 . 100 100 100 5 39 ' 48 42 44 3 9 .6 44 41 50 46 48 49 4g 6 20 26 26.9 29 24.8 28 18 23 23 25 27 27 7 11 12 15,4 18 8 9.6 14 10 14 9 7=1 11 7.1 5.8 11.7 14 4.0 5,9 14 15 7 .2 7.2 10 5,6 8 .7 5.2 5,3 7,7 9,5 3,1 . 4.7 9.1 10 5.6 5.6 11 3,0 5 .8 4 .0 4.3 4.6 ; 5 .6 2.4 3.6 4.7 5 .3 4 .2 4.2 12 2 .5 3,9 2.4 2 .6 3.6 4.5 1.5 2.3 3.3 3 .7 3 .8 3.8 13 1.9 3«0 2.2 . 2.4 14 0.08 1.3 2.7 3.0 2.1 15 0.06 1.1 S Table 82. (Continued)

NGG 6833 NGC 6884 NGC 6886 n lo Xc Io \ Xo Ic 4 100 100 100 100 100 100 5 34 44 22 . 3 46 40 51 6 20 29 9.8 26 17 25 1 11 18 9 5.4 8 .9 4.2 17 6 10 10 4.1 6 .9 2.0 8,5 2 3 .4 11 2.6 4.4 1.2 6.6 1.7 2 .9 12 , 1.7 2.9 1.1 4.9 1.9 - 3 .3

Table 8 3 . Balmer Line Intensity Bata for Nebulae Published by O'Dell (1963)

2371-2 3621 TlfnF^^iTIWT I0

Table 84. Balmer Line Intensities for J 320 Published by Burgess (1958) APPEEDXX B

Balmer Line Intensity Lata as Originally Published for Stars

219 Table 85 * BaXmer Line Intensities for y Oassp Published by Wellman (1952) n 9/26/37 11/28/37 1/22/38 2/1/38 11/28/38 11/28/38 12/12/38 A-fBC + DE + G

J1 2 3 .2 3 {IIO105*05 .O 9 0 .1 9 8 .0 2 7 .8 24.3 2 0 .0 14,60 / 19.2 16.29 8.8 9.9 6.02 4 13.80 { S i o IU 115.9 . 1 16-29 15-8 {il:!! 1 1 -9 ' 5 .1 8 r 6 ,9 2 4.66 { h f a 3.14 2.38 . 5-42 I 6 ,7 8 4.83 4.37 3.79 ' 2 .9 0 f 3 .2 5 2 .2 8 2.74 2.04 1.28 6 , {3.46 1:46 1{ 3U .4 7 r 2-28 I-98 { 1 3 f 2 . 05 7 1.51 1.51 1{ 2%% .0 8 I 1.56.5 6 { °;f| 1.59 1.59 1.03 2,18 2 .9 6 2.32 (2 .61) 1.63 1.12 1.15 O .63 AB G D E F G Equivalent Widths Central Intensities 220 821 fable 860 Balmer Line Intensities f©r Be Stars s Published by Burbidge and Burbidge (1953aj>b) and Herbig (i960)

Burbidge and Burbidge (1953a) n t» OMa 58 Bri 48 Per 0 Pse 105 fan

4 1,13 1.98 1.32 1-79

Burbidge and Burbidge (1953b) Herbig (I960) B 12 Aur 11 Can* ID 37998 . HD _ >9 1,63 1.61 2 .3I 4 ,5 1 600 0 .6 0 2 .0 7 100 I 0,-38 1.02 50 0.92 38 0,61 ©.15 0.71 08 * 9 0.76 0.20 0vl8 0-53 10 0,68 0,18 0 .5 2 ^blend 11 0 ,6 9 0.22 0.15 0-59 12 0,72 0,15 0.13 0.50 13 0.86 0.20 14 0.63 0.17 0,12 0.44 0.76 0,14 0 .1 8 0.43 0.20 0.14 S:S 17 0,53 O.lB 0.16 0 .6 7 18 O .3 4 0.15 0,15 0.40 19 0.26 0.19 0.20 0.22 20 0 ,2 5 0.13 0.17 21 0.21 0 .0 5 22 0 .1 8 0.07 23 0.10 24 0.14 25 0 ,0 6 asa Table 8 7 , Balmer Line Table 88, Balmer Line Intensities for £ Gephei^ Intensities for Nova RT Published by Wilson and Serp 1909P Published by Seddon (1956) Grandjean (1952)

n W'(x) Std. Error n I (1950) 3 1,24 ± 0.13 3 50 4 0,35 db 0,08 4 12 5 0,13 ± 0.07 5 15 6 10 7 8 8 6 9 3 10 . 1 11 1 12 1 13 1

Table 89 , Balmer Line Table 90, Balmer Line Intensities for lota Sou 1949,? Int ©ns ities for MaoRae Published by Oolaeevieh + 4 3 ° Published by (1 9 5 0) Sreemstein (1 9 5 ^)

n I (8/6-8/14/49) W ‘(X) 5/2 8/2 :/9/20/53 3 20 4 15 5 8 2 ,6 5 3.71 8,64 6 6 2,19 2.60 7 2,03 8 3 1,49 9 3 10 2 11 2 12 2 1 3 ... . 1 Table 91° Baliner Line Intensities for Hova Aqu 1918, Published by McLaughlin (1953)

n *19-*20 *21 *22 *23 ’25 ’34 ’37 00 »4l ’42 ’4? ’49 ’50

4 Tr Of . 0,5 2 1 ■ p* p 0,5 1.5 P 2 2+ 5' Tr Of 1 2-f 2+ P p 1,5 2-f P 4 4

6 ’•Of 1,5 ■ ■ p P p Of 1+ P 3 3

7 • - : - - Of P 1 *p = present and measured, but no estimate made

Table 92° Balmer Line Intensities for Nova DQ. Hereulis 193^ Published by Swings and Jose (1952) and Bajenov (1956)

and Jose (1952)

n 1940 1942'9. . 1947 ' V 19491 7 1950 3* 15 ■f + 50 100 4 8 12 12 10 10 5 6 8 : 8 10 10 6* 7 10 IGn 2 On 25 7 7 14 4 8 2 Q-» . 6 3 3 7 4 9 3 1-2 ■ . 5 1 10 2 1 3 0 11 2 2 0 223 12* 6 3 3 ___ J5 2 "^bTeiHs 224

Table 92. (Continued)

Bajenov (1956) 1934 1935 n y2/£§ 1^" T7Tr~^7r™^71T~™E^ 37m— 471 3* 2.864 2 .5 8 2 6.713 4.584 4 0,407 1 .2 2 1 I .806 6 .0 9 8 2 .0 9 8 3.245 I .798 7.684 5 °4 l6 6 .2 3 4 8 .0 6 6 7.260 6 .1 5 0 6 .1 9 0 7.824 5 - - . 4 .6 8 2 I .870 2.302 0.324 3 .9 0 8 3 ,6 9 1 4.334 4 .1 8 7 5.367 5.198 4 .5 2 5. 4.000 6* 0.287 0.556 2.118 3.834 1.277 1.011 0.843 2.492 1.820 2.210 2 ,6 6 8 2 .8 6 0 2.138 2.804 2.635 . 1 7 0 ,1 1 9 0,409 . 1 .2 6 9 3 .0 7 2 1 .1 6 7 0,.745 0.448 1.711 1.400 1 .8 0 0 1 .9 4 0 2 .2 8 9 1.281 2.404 1.474

8* O.O99 0.426 ” 0.614 0 .2 9 4 0.573 0.273 0.684 O .751 0 .8 0 0 0 .8 9 7 0.952 0.636 0.821 9 - ; ~ ^ 0 .7 9 0 0 .5 1 4 0.440 0 .3 3 2 0 .3 7 6 0 .9 3 4 o.« 0.642 0.637 0.812 0.954 •^blends

Table 93® Balmer Line Intensities for DK Lacertae 1950 Published by Wellman (1951)

1950 . 1/29 2/5 2 /1 7 3 /1 3/18 4/8 5/2 6 /3 0 7/3 4 15 : 20 15 15 20 20 15 ' 15 15 5 io' 20 ' 20 20 10 10 10 6 10 20 15 10 10 10 8 8 15* 7 8 15 10 8 5 5 4 3 3 8 3 15 8 6 9 3 10 4 4 10 2 5 1 11 1 1 12 1 13 1 225 Table $4. Balmer Line Intensities for RS Oph 1958 Published by Griffin and Thackery (1958)

. 1958 n 7/19 7/22 7/25 7/28 7 /3 0 8/10 8/11 8 /1 9 8/24 8 /2 9

3 600 600 500 600 300 400 400 300 300 250 4 150 150 100 150 80 120 100 100 100

5 60 70 50 50 50 80 4o 4o 33

Table 95• Balmer Line Intensities for T Cor Bor Published by McLaughlin (1953)

1946 1947 1949 1950 1951 1952 n 8/23 9/17 10/2 7/29 6/2 4/16 5/8 5/4

4 18 14 9 7 12 18 15 12 5 12 8 8 5 10 12 12 10 6 10 5 5 3 10 12 8 4 7 3 ■ 1.5 2 > 3 6 6 0 226 Table 96. Balmer Line Intensities Published for Two P Cygni Stars P Oyg HD 51585

"Burbidge and n Beals (1955) Burbidge (1955) Arkhipova (1 9 6 2) 3 3660 46.0 88 4 783 11.5 1 8 .5 5 325 4.38 4 .7 6 205 2 .0 8 1 .6 7 129 1.22 2.1 8 248 ------9 (12) 0.52 10 (11) 0.44

H ( 9 ) , 0.39 12 (9) 0.29 13 (9) 0.19 14 (7 ) 0.17 15 (7) 0 .1 8 16 (3 ) ~ 17 0 .0 8 18 0 .1 1

19 • 0.06 ; 20 O.O8 21 0 .0 6 22 0.05 23 0.02 24 0.02 Table 97» Balmer Line Intensities for x Oph Published by Burbidge and Burbidge (1955) and Kupo (1959)

Burbidge and Kupo (1959) Burbidge (1955) ID 24359„.» ' " n 00062 00063 *o,64 ,0.8 0 00082 00086 00.95 0 ..9 6 3 33-9 4 5-79 4 5-5 5 o2 5.3 5.3 5 .6 5 -8 5 .2 4.4 5 1 .5 8 5 2 .0 1 .6 1.5 1-7 1-7 0 .8 1.4 1 .2 6 1.03 6 0 .8 6.6 1 .1 0 .6 0 .7 0,4 2.4 0 .7 7 0.63 8 -7 0 .8 0.5 0 .5 . 9 0.44 10 0.44 11 0.41 12 0 .3 6 E p o (1959) 13 . 0.41 ' "' IB 24360 . : 14 0.44 A o..05 0-.O6 o.o07 ,o.0 8 ,09 ,..15 .. .18 ,. .19 15 0.44 5-2 T T o ™ 5 T T “ 16 0.37 0.6 0.9 0 .9 1.2 1 .5 1 .6 1.6 ' 17 0.41 0-3 l.l 1.0 18 0 .3 6 0,2 :: 0.4 0.4 1 .0 0.5 19 0.26 20 0.21 21 0.29 ^ p o (1959) n ,21 ;22 ...24 ,..40 ...43 00.44 ...46 ..o47 >53 3-9 T . y 375 TJ 3-7, 3-7 "578" 4-t /5-1 1.2 1,4 1-3: 0 .6 1.6 1.5 1.8 0 .5 0 .5 0.6 5 3 0.4 0.7 227 Table 98» Balmer Line Intensities for BP Oyg Published by Merrill (1950) and Aller (1955)

Merrill (1950) Aller (1955)

^q a r '-alq 1950 :1951 n 1940-4 % 5 /2 1 6 /2 0 7/9 7/15 5/17 9/7 1 0 /1 6 17/19 7 /2 0 7 7 /2 7 17/27

3 300 4 60 218 820 550 125 520 5 30 56 270 38 161 36 170 6 18 20 100 27 111 21 63 84 4? 100 7 12 33 39 22 35 19 44 63 42 52 8 10 6.5 26 12 15 11 35 20 28 ¥ 9 8 5 11 3.0 3.4 13-2 l4.4 9-5 15-9 10 8 8,8 10.3 7-2 12.9 11 7 7-1 8.3 6,0 10.2 12 . 7 - 8.2 5-2 5-0 8.4 13 6 5.0 4.6 3-0 5-1 14 6 7-4 4.6 4.0 6 .9 15 5 2,8 5-4 3-3 6,2 16 5 3-9 3-4 3-0 6,1 17 5 1.4 3-5 2.1 4.2 18 4 2.2 3-5 1.2 2.3 19 4 1.0 1.6 20 4 0.5 0.92 21 3 22 3 - I W«(X) 23 3 24 2 25- t • 28 1 1 ' 4 855 229 Table 99o Balmer Line Intensities for 01 Oygnl Published by Merrill (1950) and Aller (1955)

Merrill (1950) Aller (1955) 1950 ii 1948-49 1951 7/16 7/16 8/12 9/19 3 300 4 • 60 2200 100 50 50 5 30 430 50 25 30 6 20 3.60 35 30 12 Z 15 12 10 5 8 12 5 3 1 9 T 10 7 11 6 W ’(X) 12 . 5 14 4 15 4 lo 4 17 3 18 3 19 3 20 3 21 2 22 2 23 2 24 2 25 1 26 1 27 1 28 1 29 tr 30 tr 31 tr 32 tr 33 tr 230

Table 100. Balmer Line Intensities for y Cygni Published by Mao-Lin (1950) and Fugita (1955)

Mao-Lin (1950) F^glta (195*) n 7/48 10/8 10/10 11/6 11/23 12/12 3 30 4 25 1.62 1.67 2.01 1 1 .0 8 3.53 5 20 1.55 1 -5 0 3.32 5 .6 6 .6.14 6 15 3.79 3.76 3.97 112.05 7.00 T weak weak --weak . 0.12 8 11 6.30 5.57 5.52 3.08

10 ro 4.00 1.56 ■;

9 .8 10 ; 8 0 .5 4 1.66 11 9 0.38 v ■ 1.03 12 5 • - ' ill

13 ; 2 ' - r ' S 14 15 : 4 16 17 8 18 4 II 19 T a b le 1 0 1 . Balmer Line Intensities for AG Pegasi Published by Mao-Lin195 ( 0a ) , Mao-Lin and Bloch1 9 ( 5 2),; Burbidge and Burbidge 1 ( 9 5 4) , and Arkhipova and Dokuchaeva (1962)

Mao-Lin and Burbidge and Arkhipova and Burbidge fiQPd n 1 9 4 S Bloch (1952) Dokuchaeva (1962) 1948 17 300 >300 >300 I 16 165 : 160 140 188 160 2 3 .8 1 7 .9 3 6 .8 14 100 .. 100 100 60 100 29 1 5 .2 1 6 .4 18.6 I 13 65 . 66 70 56 87 18 ■1 3 .5 9 .2 13-7 12 42 42 65 33.4 58 4 .5 6 .9 3.3 9.3 I 11 50 50 60 36.9 (68) ■ 4.6 4 .6 7.1 9 7 35 35 40 14.87 27 1.6 2 .4 4.1 10 8 IO .65 22 0 .9 4 I I 6 10.18 21 1 .6 4 12 6.45 14 1.4 1 4.68 10 11 3< 3.48 7.6 3 4.10 9.0 i! 1. 2.55 5.6 2 1 .8 3 4.1 II 3 1.81 4.1 19 1 .6 9 3.8 20 1.35 3.1 21 1.48 3.3 22 1.18 2.7 1.04 2.4 II 1.02 2.3 0.91 2.1 II 0.91 2.1 0.46 1.1 11 0.67 1.5 29 0.64 1.5 -29 0.57 _ J L 2 . Decrement 232 Table 102„ Balmer Line Intensities for MWC 603* Published by'Tlfft and Greensteln (1958)

n 3 4 22 25 5 18 20 6 15 17 7 12 14 8 10 11 9 9 8 10 8 7 11 7 5 12 6 4 13 5 3 14 4 3 15 4 3 16 3 2 17 3 2 18 3 2 19 3 20 3 21 2 22 2 23 2 24 2 25 2 26 1 27 1 APPENDIX C

Balmer Line Intensity Data as Originally Published for the Solar Chromosphere

233 234 Table 103. Balraer Line Intensities for the Solar Chromosphere as Given In Some Early Publications

Schwartzs child Russell ^gtratton^ Mitchell Mitchell (1906) (1926) (1926) (1930) (194?) 3 10 200 200 4 9 200 160 5 30 8 160 150 6 20 7 40 140 14 7 10 5 35 120 4 8 8 4o 100 140 9 8 30 90 120 10 8 30 80 100 11 7 25 70 80 12 6 , 25 - 70 75 13 6 20 55 70 14 5 20 50 60 15 4 15 45 50 16 4 15 40 45 17 3 15 35 40 18 2 15 30 40 19 3 15 25 35 20 2 15 25 30 21 2 15 20 30 22 2 11 18 25* 1 23 2 9 18 22 24 9 18 20+ 25 7 12 22 26 7 12 15 27 5 15 12 28 5 7 15-S- 29 6 15 10+ 30 4 , 3 18+ 31 6 3 6 32 3 2 4+ 33 1 3 3 34 5 2 2 35 1 2 3+ 36 1 4 37 0 2 Table 104. Logarithms of the Intensities of the Balmer Lines in the Solar Chromosphere as Published in Clllid and Menzel (1936)

n h 670 .1500 2300 3170 '4000 3 16.46* 15.08 14.59 4 15.76t 15.54 15.11 14.68 13.49 14.18 13.98 5 15.17 14.97 14.55 14.01 13-58 6 14.83 14.72 14.26 1 3 .6 9 13-46 1 3 .0 5 7 14.61 14.57 14.15 12.94 8 14.42 14.53 14.01 12.81 12.94 9 14.35 14.31 13.91 12.69 10 14.11 14.09 13.67 13.28 12.65 11 14.05 < 13.61 13.17 12 .5 2 12 14.02 13.58 13.11 .13 13.98 13.53 13-04 14 1 3 .8 8 13.43 1 3 .0 0 15 1 3 .8 0 13.33 1 2 .8 5 16 13.74 13.25 12.84 17 13.69 13.23 1 2 .7 9 18 1 3 .6 0 13.17 19 13.58 13.11 20 13.51 13.01 21 13.50 12.99 22 13.40 1 2 .9 2 , 23 13.40 12.87 24 .13.31 1 2 .7 2 25 I3 .2 8 12.71 26 13.26 27 13-20 28 13.17 29 13.10

30 13.08 235 31 _ 13.05 ...... ’^Grating spectrograph iXJV spectrograph 236 Table 105» Logarithms of the Intensities of the Balmer Lines In an Exeited Region of the Solar Chromosphere as Published in Oillid and Menzel (1936) n h 900 1730 2570 70 900 1730 2570

I 15.90 14.84 14.18 15.31 15.00 14.45 13.70 6 14.97 14.79 14.15 7 14.55 13.94 8 15.05 14.66 13.99 9 14.49 14.85 14.39 10 14.25 12.67 14.76 14.26 11 14.19 12.56 14.72 12 14.16 13.46 1 2 .5 2 13 14.12 13.41 12.55 14 14.02 13.28 12.45 15 13.94 13.23 16 13.88 13.13 17 13.83 13.11 18 13.74 12.97 19 1 3 .7 2 12.97 20 13.65 12.84 21 13.67 12.86 22 13.57 12.74 23 13.52 12.65 24 13.49 12.64 25 13.44 12.57 26 13.40 12.49 27 13.30 12.45 28 13.32 12.39 29 13.25 30 13.35 31 13.09

TV film Grating film Table 106. Logarithm of Intensities of the Balmer Lines in the Solar Chromosphere as Published in Vyazanitsyn (1951)

Second Contact n h' 3050 550 3900 1500 350 4850 2600 1900 1100 2350

3' 14.86 15.35 14.21 15.19 15.33 14.98 15.31 1 5 .3 6 14.77 4 13:73 15.09 13.27 14.93 15.18 13.18 14.59 14.89 1 5 .1 5 14.34 . 5 13.30 14.84 1 2 .9 0 14.80 15.02 12.94 14.33 14.65 1 5 .0 0 14.02 6 13.21 14.72 1 2 .8 0 14.63 14.89 12.81 14.07 14.50 14.82 1 3 .8 8 7 13.16 14.61 12.65 14.45 14.78 12.60 13.94 14.36 14.68 1 3 .7 8 8 13.12 14.46 12.54 14.33 14.66 12.37 13.79 14,20 14.56 1 3 .7 1 9 12.70 . 14.16 12.55 14.15 14.33 12.07 13.11 14.17 14.22 1 3 .6 8 10 12.36 14.07 I2 .3O 14.02 14.16 11.96 12.78 1 3 .9 5 14.45 1 3 .4 7 11 1 3 .9 2 13.76 14.07 1 2 .8 2 1 3 .5 6 14.03 1 3 .2 2 12 13.80 13.73 14.08 12.75 13.50 14.01 1 3 .1 5 13 13.81 13.73 14.09 12.57 13.47 13.95 13.14 14 13.71 13.48 1 3 .8 6 12.54 13.28 13.75 12.94 15 13.52 13.28 13.76 12.46 13.12 13.76 12.76 16 13.4-6 13.23 1 3 .7 2 12.46 12.95 13.61 12.60 17 13.14 12.92 1 3 .2 7 12.42 12.48 13.30 12.31 v 18 13.06 1 2 .8 5 12.57 12.21 12.84 11.72 point 1 point 2 point 3 237 Table 106 (Cont„). Logarithm of Intensities of the Balmer Lines in the Solar Chromosphere as Published in Vyazanitsyn. (1951)

' n 5700 3600 2950 2150 3250 7650 575O ~ 5200 4350 5200 3 14.84 14.84 14.89 14.60 14.22 4 1 2 .6 0 1 5 .8 0 14.64 14.80 14.22 1 3 .0 6 13.36 13.43 12.75 5 1 2 .1 0 13.39 14.22 14.25- 13.96 1 2 .6 6 12.95 1 3 .0 0 12.44 6:, 1 3 .0 6 13.91 14.22 13.76 12.46 12.64 1 2 .8 0 1 2 .2 0 7 1 3 .0 1 13.67 14.15 13.56 12.32 12.46 1 2 .6 2 12.05 8 13*44 14.10 13.40 1 2 .2 2 12.49 1 1 .9 6 9 13.14 1 3 .6 9 13.30 12.39 10 13.00 1 3 .5 0 13.18 12.32 11 1 2 .8 5 13.43 1 2 .7 0 12 12.77 13.30 12.65 13 12.76 13.18 1 2 .6 5 14 12.55 12.99 12.48 15 12.41 1 2 .7 6 12.36 16 12.37 1 2 .6 3 . 12.04 17 11.74 12.10 11.95 18 11.86 point 4 point 5 Tabl.6 106 (Cont.). Logarithm of Intensities of the Balmer Lines in the Solar Chromosphere as Published in Vyazanitsyn (1951) n b 8600 ' 6900 6300 5450 6250 8050 7500 6650 7300

3 13.85 x/ 4 12,14 12.54 1 2 .8 8 1 2 .9 0 11.93 12.35 12.47 5 1 1 ,6 2 12.13 12.48 1 2 .3 6 1 2 .0 6 1 2 .0 0 6 II .9 0 12.25 1 2 ,2 2 11.88 11.60 7 11 .78 12.11 1 2 .0 6 11.76 11,34

• "point 9 point 10 239 Table 106 (Cont.). Logarithm of Intensities of the Balmer Lines in the Solar Chromosphere as Published in Vyazanitsyn (1951)

Third Contact

XL* 8200 6950 6250 7700 5750 4900 6700 4600 3550 5850 3600 2400 3 '13.47 13.94 14.48 12.64 14.86 1 5 .1 0 1 3 .1 0 1 5 .0 0 4 1 2 .1 0 12.35 12.05 12.79 12.47 11.91 13.36 14.44 12.25 14.52 14.89 5 1 1 .8 2 1 1 .3 0 1 2 .3 0 12.15 11.27 1 2 .7 6 13.83 1 3 .2 2 14.33 6 11.48 11.15 1 2 .1 0 11.75 12.46 13.57 1 2 .8 3 13.99 T 11.24 11.95 11.49 1 2 .2 6 13.31 1 2 ,5 6 13.75 8 11.79 1 2 .1 0 1 3 .1 0 1 2 .3 8 13.61 9 11.93 12.67 1 2 .3 0 13 .2 1 10 11.75 12.44 12.02 1 2 .8 7 11 11.36 1 2 .3 0 11.99 12.77 12 12.15 11.84 12.76 13 11.92 ' 11.70 12,69 14 11.49 12.39 15 U . 4 o 12.22 16 11.38 12.29 17 11.06 11.52 18 10.86 point 11 point 12 point 13 point 14 point 15 OtiS Table 106 (Oont.). Logarithm of Intensities of the Balmer Lines in the Solar - Chromosphere as Published in Vyazanitsyn (1951) n 5100 '2600 1300 1800 - 4150 1550 100 3200 400 3 13.27 15.12 15.66 13.96 15.73 15.25 1 5 .4 4 k 12,37 14.61 15.25 13.81 15.38 15.50 14.44 1 5 .1 8 5 11.97 14.00 15.09 14.64 13.10 14.92 14,90 12.78 14.88 6 11.73 13.55 14.75 14.41 12.67 14.38 14.57 14.68 T 11.55 13.30 14.53 14.23 12.41 14.20 14.35 14.52 8 13.13 14.39 14.03 12.18 14.10 14.19 1 4 .4 l 9 13.03 14.10 13.50 13.50 13.81 14.05 10 12.65 13.04 . 13.18 13.82 11 12.45 12.93 13.09 1 3 .5 4 12 12.18 1 2 .9 8 13.05 14.49 13 11.94' 12.93 13.39 14 - 11.71 12.63 13.26 15 11.49 12.39 12.98 16 11.53 12.25 12.84 17 11.50 11.46 18 11.21 point 16 - point 17 point 18

ro 4^ H 242 Table 107» Logarithms of the Intensities of the Balmer Lines in the Solar Chromosphere as Published in Vyazanitsyn (1952)

h 650 1000 1750 2150 3*50 3 15.73 15.41 1 5 .5 4 15.26 14.83 4 1 5 .1 2 15.07 14.91 14.80 14.38 5 14.95 14.99 14.79 14.00 IS.?? 6 14.84 14.77 14.46 14.06 13.54 8 14.33 14.02 14.2? 13.92 13.11 9 14.18 13.70 13.84 13.15 12.84 10 14.06 13.15 13.53 13.33 11 13.15 13.34 13.28 12.88 12 13.57 12.88. 12.75 12.59 point ]L point 2 point 3

n h 4250 4950 5650 6200 4550 2150 600 3 14.49 14.09 15.04 15.50 4 13.86 13.48 1 3 .6 9 13.20 13.61 14.40 14.84 5 13.06 13.07 1 2 .9 8 12.72 12.93 14.05 14.55 6 12.69 12.90 12.53 13.81 14.46 ' 8 12.82 1 3 .9 6 9 13.61 13.81 point 4 point 5 point 6 point 7 243 Table 108. Logarithms of the Intensities of the Balmer Lines in the Solar Chromosphere as Published in Athay, Menzel, and Orrall (1957)

h 1000 1720 2440 3140 n 290 3850 4550 5250 5950 5 4.77 4.57 4.16 3.63 3.39 2.77 2 .2 5 6 4.86 4.47 4.12 3.72 3.22 3.09 2 .7 0 2.27 8 4 .9 6 4.66 4.16 3.76 3 .3 0 2 .8 7 2.62 9 4 .6 9 4.29 3.68 3 .1 6 2.84 2 .58 10 4.54 3.53 2.94 2 .6 7 11 4.51 3.35 2.72 12 4.47 3.27 2.71 13 4.43 3.25 2.59 14 4.36 3.17 15 4.13 2 .9 7 16 4.06 2 .8 3 17 3.97 18 3 .8 8 19 3.87 20 3.76 21 3.88 22 3.80 23 3.65 24 3 .6 2 25 3.63 26 3.38 27 3.35

1 ^ 5 6 0 1360 2140 2900 3700 4480 5220 6060 6800 7600 inii-iiii m ini < n fLp . inr ' T W - , l . j . 1:!- ■-;,iniiifi'woi iii;riiiinin irimiMir inLLJ hh u.i.ju im.iUm.iMi humi-l.i.Hi u , ir.i. m _ .1. ijjim n 1 ujn n n i 1, r.iuiljj iiiu wumiii .... . 11.11M .1111 w ,T . j mlTi.imh ni j u n , .. . 11#,i n „ n , lm i.l n < nrjr.u.^ninu 3 ' . 5.34 5.14 4.76 4.24 3 .9 6 4 5 .5 2 5.24 4.93 4.75 4.41 3.97 3.35 2.96 2 .6 6 2 .0 7 Table 109» Logarithms of Intensities of the Balmer Lines for the Solar Ohromosphere as Published in Houtgast (1957)

h 1000 564 n 129 553 1605 132 1020 1638 4 3.27 3.32 2.86 2.73 5 3.24 3.08 2.59 2.23 6 2.23 1.95 7 3*63. 2.56 2.40 2.71 2.54 2.07 1.58 . 8 3.24 3 .1 1 2.88 2.62 2.77 2.47 2.16 1.75 . 9 3.02 2.95 2.65 2.16 2.41 2.03 1.70 1.17 10 2 .83 . 2 .53 2 .3 8 I .8 7 2 .1 1 1 .9 6 1.46 1.09 11 2.75 2 .5 0 2.27 1.79 2.09 1.79 1.30 1 .0 2 ' strip a strip b

h h 129 553 125 536 970 1557

4 3.23 3 .1 5 2 .8 i; 2.72 2.57 2.09 5 2.99 2.91 2.41 2.29 I .98 1 .6 0 6 2 .7 6 2 .6 8 2.04 1 .8 0 1 .6 2 1 : 18 7 2.41 2 .2 1 1.64 1.45 1 .2 2 0 .8 6 8 2.47 2 .2 1 1 .8 0 . 1 .5 8 1.46 1 .1 2 9 2 .0 3 1.75 10 1.74 11 1.84 strip c strip d

Table 110. Logarithms of Intensities of the Balmer Lines for the Solar Chromosphere as Published in Houtgast (1962)

-4-1000 o n o o o "2000 :3 0 0 0 .4000 ;.50oo 6000 77000 8000 1% Table 111, Logarithms of Intensities of the Balmer Lines in the Solar Chromosphere as Published in Thomas and Athay (1961)

h 100 640 750 960 1070 1180 1280 1390 n 530 850 ”” 3 ' . 5776"~’' . k 5.68 5-75 5.62 5.69 5 .6 2 5.55 5.66 5.52 5.52 5 5 .6 6 5.68 5.62 5 .6 3 5 .5 2 5.50 5.56 5.44 5.37 6 5.51 5.55 5.47 5.42 5.35 5.36 5.21 5.11 8* 5 .2 7 5 .2 7 5.07 5.12 5 .0 8 5.03 5 .0 0 4.84 4 .8 5 4.79 9 5.01 4.88 4.86 4.80 4.73 4 .7 2 4.62 4.58 4.51 10 5.04 4 .9 0 4.77 4,75 . 4.69 4.63 4.59 4.50 4.44 4.4o 11 4.99 4,80 4.70 4.67 4.6o 4.55 4 .5 1 4.4o 4.36 4.28 121 4.97 4,80 4.68 4.66 4.59 4 .5 1 4.46 4.36 4.30 4.18 13i 4 .9 8 4 .7 8 4.64 4.62 4.54 4.48 4.46 4.31 4.22 4.14 Ik 4,86 4,71 4 .5 8 4.60 4.45 4.41 4.34 4.22 4.07 4.06 15 4,69 4 .5 8 4.44 4.41. 4.29 4.25 4.18 4.07 3.99 3.91 16 4,64 4.49 4 .3 6 4.39 4.20 4.22 4.08 3.98 3.97 3.83 IT 4.57 4.40 4 .28 4.30 : 4 .1 7 4.13 4.03 3.96 3.85 3.76 18 4.43 4.35 4.22 4.20 4 .0 7 4.01 3.93 3.90 3.76 3.66 191 4.45 4.31 4.20 4.19 4 .1 1 3.98 3.92 3.84 3.73 3.69 20 4.42 4.26 4.17 4.13 4 .0 1 3.90 3.78 3.76 3 .6 9 3.59 21 4.37 4.27 4.17 4.15 4 .0 1 3.77 3.75 3.75 3.61 . 3.51 22 4.35 4.12 4.02 3.99 3 .9 1 3.68 3.68 3.65 3.60 3.47 23 4.27 4.08 3.97 3.94 3 .8 0 3.59 3.57 3.52 3.40 3.35 24 4.25 4.05 3.93 3.88 3 .7 8 3.56 3.51 3.45 3.32 3.24 25 i 4.26 4.05 3.89 3.86 3 .7 5 3.46 3.47 3.43 3.32 3.19 26 3.90 3.81 3.76 3 .6 9 3.39 3.39 3.34 3.26 3.16 27 4,18 3 .8 7 3.78 3.68 3 .6 5 3.43 3.35 3.21 3.15 3.10 28 f 4.29 3.94 3.87 3.70 3 .7 1 3.34 3.35 3.16 3.16 3.04 29 3.73 3.67 3.60 3 .5 4 3-18 3.26 3 . 1 1 1 3.00 . 2.88 30* 4,05 3.88 3 >80 3 .7 1 3.32 • 3.36 3.17 3.21 2.98 31 3.57 3.58 3.48 3.07 3.14 2.94 2.82 2,73 ^Strong blends Weak blends Table 111 (Cont„). Logarithms of the Intensities of the Balmer Lines in the Solar as Published in Thomas and Athay (1961)

1730 n 1500 1610 2520 TW TW 4 5 .4 5 5.36 5.38 5.32 5.22 5.15 5 .3 1 5 .2 7 5.20 5 .1 2 4 .9 8 4.78 I 5 .0s 5 .0 4 5.04 4.70 4.49 8^ 4=80 4 .7 4 4.63 4.63 4.55 4.46 4.44 4.38 4.44 4.09 9 %.53 4.46 4.25 4.47 4.20 4,12 3.96 3.92 3.33 3.70 10 4 .3 6 4.32 4.12 4 .1 5 3.97 3*95 3.85 3 .8 0 3.73 3.55 II 4.28 4.22 4.00 4,02 3.97 3.88 3.78 3.64 3.60 121 4,1? 4.13 3.91 3.88 3 .8 1 3.70 3.63 .3-66 3.52 3:8 131 4.10 4.03 3.80 3.74 3.79 3.82 3.50 3.48 3.49 3.25 14 • 3.95 3.97 3.72 3.62 3,61 3.69 3.57 3 .4 5 3.40 3-20 3.85 3.85 3.63 3.48 3.56 3.48 3.30 3.19 3.27 3.10 $ 3.80 3.77 3.94 3.39 3.50 3.38 3.20 3.09 3.07 17 3.67 3.65 3,47 3.32 3.46 3.34 3.04 3.06 3 .0 8 2.92 18 3.65 3.50 3.42 3.21 3.33 3 .1 6 2.94 2 .9 7 2.84 2 .7 8 19t 3.61 3.60 3.32 3 .1 3 3.33 3 .1 5 2.95 2,88 2.88 ■2.66 20 3.59 3.40 3.23 3.33 3.21 2.99 2.88 2.78 .2.77 21 3 .5 0 3.34 3.15 3.20 3 .1 1 2 .9 0 ■ 2.73 2.74 2.72 22 3.33 3.23 3.11 3.18 2.86 2 .7 0 2.64 2.61 23 3.22 3 .1 7 3.04 3.02 2.86 2 .7 8 2 .7 0 2.62 2.62 24 3.19 3.04 2.93 2.93 2.81 2.69 2 .5 8 2.55 2.50 25t 3.10 - 2.97 2 .9 0 2.77 2.76 2 .6 2 2 .5 6 2.47 2.27 26 3.03 2.94 2.85 2.73 2 .7 2 2.59 2 .5 2 2.35 2.39 27 2 .9 6 2.83 2 .7 6 2.64 2.65 2.52 2.35 2.38 2.23 2 8 i 2.91 2 .8 1 2.79 2.63 . 2 .6 0 2.43 2.42 2 .3 0 2.36 29 2.74 2.73 2.76 2.45 2.49 2.39 2.28 2.31 2.20 30* 2 ,80 2.82 2.75 2 .5 2 2 .6 0 . 2 ,4 5 . 2 .3 8 2.35 2.27 31 2.63 2.58 2.60 2.36 2.37 2 .2 6 2.22 2.14 _"!ETen3s" ro fWeak blends Table 111. Logarithms of the Intensities of the Balmer Lines in the Solar Chromosphere as Published in Thomas and Athay (1961)

h 3180 3720 3870 4680 5490 6300 r n r w ” i r r ~ T ( F ~ n 3060 2000km. Max. h t . 3 5.44 3 T 8 F 0.35 — T 3 ( r 4 4.74 4.25 3,08 2.64 0 .8 7 1.60 5 4.54 3.94 3.23 2.53 1.20 1 .8 0 6 4.20 3*60 2«9l 2.42 1.36 1 .8 0 8* 3 .7 4 3.79 3.44 1.39 1.39 9 3 .3 8 3.30 3 .0 8 1.51 1.51 10 3.24 3.22 2 .9 2 1.56 I .56 11 3.12 3.00 2 .8 2 1.62 1.62 127 2.97 2 .9 0 2.55 1.68 1.68 137 2.88 2.84 1.68 1.68 14 1.70 1 .7 0 1.73 1.73 il 1.73 1.73 17 1.79 1.79 18 I .85 1.85 19t 1.91 1.91 20 1.91 1.91 21 1.96 1.96 22 •1.96 1.96 23 "1.96 1.96 24 2 .0 5 2.05 257 2.08 2.08 26 2 . % 2.C5 27 2 .0 8 2 .0 8 2.12 2.12 . 29 2.08 2.08 30* ' 2.12 2.12 31 2 . 1 2 2.12 "^§trong™bIeMs" W e a k . blends 248 ^able 112. Logarithms of Intensities of the Balmer Lines in an Exoited Region of the Solar Chromosphere as Published in Thomas and Athay (1961)

IT ■ 400 1350 n 290 510 620 810 1620 2160 4590 3 6 .1 2 5 .9 8 6 .1 1 " 5 .9 6 5 .7 6 5.33 4 6.01 5.97 6.04 5 .8 0 5 .5 2 5.19 5 5 .8 7 5.78 5 .8 3 5-64 5.40 5.04 8 5 .6 0 5 .1 8 5 .2 5 5.11 5 .1 0 9 5.39 4.98 5.01 4.83 4.64 10 5 .2 5 4.85 4.68 4.61 4.47 11 5 .0 6 4.76 4.61 4.50 4.38 12 5.07 4.73 4.55 4.45 4.32 13 5 .0 1 4.65 4.58 4.40 4.28 14 4.87 4.55 4.49 4.26 4.08 15 4.70 4.48 4.29 4.13 4.01 16 4 .5 6 ,4.38 4.19 4 .0 3 3 .9 6 17 4.46 4.27 4.13 3 .9 0 3 .8 2 18 4.40 4.19 3.9b 3.80 3.74 19 4.35 4.25 3.93 3.75 3.65 20 4.25 4.10 3.79 3.71 3 .7 0 21 ' 4.18 4.03 3.85 3.61 3.56 22 3 .9 6 3 .7 0 3.53 3.46 23 3.95 3.69 3.30 24 3 .7 7 3 .5 9 3 .2 1 25 3 .7 8 3 .5 7 3 .1 3 26 3 .7 1 3 .4 7 3 .1 1 27 3 .6 0 3 .3 2 2 .9 8 28 3 .5 4 Table 113. Logarithms of Intensities of the Balmer Lines in the Solar Chromosphere as Published in Martynov and Abluseva (1962)

h 2210 0 530 980 1480 3290 2740 n 730 1260 1710 4 1 5. oo 14.47 14.04-9 14.53 13.44 14.00 5 14.09 13.51 12.84 13.50 6 14.14 13.97 13.53 13.23 14.30 14.16 13.84 12.55 1 3 .0 7 8 13.94 13.66 13.36 12.99 14.40 14.11 13.80 13.57 . 9 13.95 13.54 13.20 12.67 14.43 14.10 13.82 13.49 12.46 10 13.60 13.23 12.80 12.36 14.02 13.72 13.49 13.14 12.18 11 13.51 . 13.11 1 2 .7 2 12.26 13.99 13.70 13-39 13.00 12.09 12 13.47 13.07 12.65 12.21 14.00 13.72 13.31 12.94 12.01 13 13.44 12.99 12.62 12.24 14.01 13.70 13.31 12.86 11.97 14 13.31 12.93 1 2 .5 0 12.11 13.84 13.55 13.18 12.76 11.96 15 13.12 1 2 .7 1 12.33 12.04 13.67 13.40 1 2 .9 8 12.62 16 13.07 12.68 12.28 13.68 13.37 1 2 .9 6 12.61 17 12.93 1 2 .5 2 12.28 13.36 13.23 12.79 12.44

ro -Pr Table 113 • (Oonfcimied)

— TT 2160 1430 760 2610 2050 1470 740 2490 1940 1360 630 n 4 1 4 .5 0 13.88 14,40 14.03 14.48 5 14.13 14.34 13.49 14.00 13.76 14.13 6 13.57 14.37 13.04 13.56 14.18 13.28 13.67 14.22 8 .13.91 14.20 12.86 13.32 1 3 .7 8 14.29 13.37 13.91 14.42 9 1 3 .7 9 14.22 12.40 1 3 .0 2 13.64 14.34 13.11 13.60 14.45 10 13.49 13.88 12.21 12.62 1 3 .3 0 13.98 12.71 13.36 14.18 11 1 3 .3 4 13.86 11.95 12.46 13.25 13.89 12.66 1 3 .2 9 14.09 12 13.34 1 3 .9 2 11.97 12.42 13.25 13.98 1 2 .5 0 13.11 14.09 13 13.30 13.98 1 2 .3 2 1 3 .2 1 14.12 12.40 13.15 14.22 14 1 3 .1 0 13.78 12.24 13.06 13.67 12.31 13.16 13.88 15 1 2 .8 9 13.60 1 2 .8 9 13.60 12.16 12.92 13.72 16: 13.03 13.62 12.94 13.52 12.15 1 2 .9 6 13.72 1? 12.69 13.41 1 2 .7 0 13.23 12.07 1 2 .7 2 13.39 250 APPENDIX D

Intermediate Steps in the Computer Computations

251 fable 114, The Data Inserted into the Computer Program n ;iog%e"in[idoo) @n(iooo) logT^graoss) igaoob) logaelnCfOOO) fTj

3 15.58 0.75 * 10 km 15.37 0 .7 5 • 10 ten 14.71 1 .0 6 4 15.19 1.20 14.86 1.20 13.85 1.40 5 15.08 1.32 14.57 1.32 1 3 .5 8 1.49 6 14.82 1.34 14.35 ,1.34 13.25 1.46 I 14.57 1.36 14.09 I .3 6 13.09 1.36 9 14.31 1.52 13.74 I .52 12.77 1.52 10 13.90 1.55 13.30 1.55 11 14.04 1.69 13.32 I .69 12 13.95 1.71 13.22 1.71 13 14.03 1.88 13.31 1.88 14 13.90 1.88 13.18 1.88 15 13766 1.80 13.14 1.80 15 13.68 1.74 12.93 1.74 17 13.63 1.88 1 2 .9 8 1,88 18 13.65 I .92 12.97 1.92 - 19 . 13.54 2.0 13.15 2.0 20 13.60 2.0 12,99 2.0 21 13.41 2.1 12.90 2.1 22 13.32 2.2 12.86 2.2 23 13.23 2.1 12.78 2.1 24 13.16 2.2 12.69 2.2 25 13.18 2.2 1 2 .6 2 2.2 26 . 13.03 2.1 12.59 2.1 27 13.13 2.1 1 2 .5 2 2.1 28 13.34 2.1 12.43 2.1 29 13.18 2.1 12.39 2.1 30 13.32 2.1 12.45 2.1 31 13.07 2.1 1 2 .2 6 2.1 sSs Table 115. The Fortran Program Used to Obtain Tables of {logj0In(h)} fi'om One Value of Pn (h0 ) and In (h0 ) for Each n

SEQ. STMH1 FORTRAN STATEMENT 1 ' DIMENSION Y(l4) 2 DO 100 1=1,3 3 READ 1,AI,B»H0 4 1 FORMAT (F6.2,F5.2,F4.1) 5 DO 101 J=l,29 6 READ 2,AIN,BN 7 - 2 FORMAT (F6.2,F5»2) 8 DO 102 K=l,14 9 C=K 10 H=0/2. 11 102 ■ Y( K)=AIN”M-» (H~H0) /2.3 12 101 PRINT 3»Y 13 3 FORMAT (14F8.2) 14 PRINT 4 15 4 FORMAT (1H1) 16 100 CONTINUE . 17 END

Table 116. The Fortran Program Used to Obtain Tables of {iTChl*} froni 0ne Value of .Pn (h0 ) and In (h0) for Each n

SEQ. STMNT FORTRAN STATEMENT 1 DIMENSION Y(l4) 2 DO 100 1=1,3 3 READ 1,AI,B,H0 4 1 F0RI4AT (F6.2,F5.2,F4.1) 5 DO 101 J=l,29 6 READ 2,AIN,EM 7 2 FORMAT (F6.2,F5.2)' 8 DO 102 K=l,14 9 . 0=K 10 H=0/2. 11 X=AIN-AI+(B-BN)*(H-H0)/2.3+2. 12 102 y(K)=10.**X 13 101 PRINT 3dY 14 3 FORMAT (mo, 14F8.1) 15 PRINT 4 16 4 - FORMAT (IHl) 17 100 CONTINUE 18 END Table 117- The Fortran Program Used to Obtain the Average {log10In (h)}

SEQ. STMNT FORTRAN STATEMENT 1 DIMENSION AI(3)»AIN(3)»B(3),BN(3)»H0(3) 2 DIMENSION Y(14) 3 DO 10 NN=15 3 4 10 READ 1,AI(NN)>B(NN).HO(NN) 5 1 FORMAT (F6.2,F5.2»F4.1) 6 DO 101 J-1,29 7 IF(j-7) 2 0 ,2 0 ,3 0 8 20 LL=3 9 GO TO 25 10 30 LL=2 11 GO TO 25 12 25 DO 103 L-l.LL 13 103 READ 2,AIN(L),BN(L) 14 2 FORMAT (F6.2,F5-2) 15 DO 102 K=l,l4 16 G=K 17 H=C/2 . 18 Z=0 19 DO 40 M=1,LL 20 40 Z=AIN(M)-BN(M)*(H-H0(M))/2.3+Z 21 BB=LL 22 102 y (k )=z /b b 23 101 PRINT 3,Y 24 3 FORMAT (14F8.2) 25 END 255 Table 118. The Fortran Program for Obtaining r ln (h)"i the Average (jj^Th'Tj

SEQ. STMT FORTRAN STATEMENT 1 DIMENSION AI(3),AIN(3),B(3)»EN(3)#H0(3) 2 DIMENSION Y(l4) DO 10 NN=1,3 I 10 READ 1SAI(NN),B(NN),H0(NN) 1 FORMAT (F6.2,F5.2,F4.1) I DO 101 1=1,29 IF(J-T) 2 0 ,2 0 ,3 0 I 20 LL=3 9 GO TO 25 10 30 LL=2 11 GO TO 25 12 25 DO 103 L=1,LL 13 103 READ 2,AIN(L),BN(L) 14 2 FORMAT (F6.2,F5.2) DO 102 K=l,l4 ii C=K 17 H=C/2. 18 Z=0 19 DO 40 M=1,LL 20 X=AIN(M)-Al{M)+(B(M)-BN(M))*(H-H0(M))/2.3+2. 21 40 Z=10.**X+Z 22 BB=LL 102 Y(K)=Z/BB 11 101 PRINT 3(Y 3 FORMAT (1H0, 14F8.1) 11 END 256 Table 1191 -Values of log10l^(h)iComputed from O n(1000km)} and {ln(iobokm)}

h 'n 500 1000 1500 2000 2500 3000 3500 3 15.T2 15.25 15.09 14.93 14". 7^ 4 15-45 1 5 .1 9 14.93 14.67 14.41 14.15 13.89 5 15.37 1 5 .0 8 14.79 14.51 14.22 13.93 13.65 6 15.11 14.82 14.53 14.24 13.95 13.65 13.36 8 14.87 14.57 14.27 1 3 .9 8 1 3 .6 8 13.39 13.09 9 14.64 14.31 13.98 1 3 .6 5 13.32 12.99 12.66 10 14.24 13.90 13.56 1 3 .2 3 12.89 12.55 12,22 11 14.41 14.04 13.67 1 3 .3 1 12.94 12.57 12.20 12 14.32 13.95 13.58 1 3 .2 1 12.83 12.46 12.09 13 14.44 14.03 13.62 1 3 .2 1 12.80 12.40 ■ 11.99 14 14.31 13.90 13.49 1 3 .0 8 12.67 15 14.05 13.66 13.27 12.88 12.49 16 14.06 13.68 13.30 1 2 .9 2 12.55 IT 14.04 13.63 13.22 12.81 12.40 18 14.07 13.65 13.23 12.82 12.40 19 13.97 13.54 13.11 12.67 12.24 20 14.03 13.60 13.17 12.73 12.30 21 13.87 13.41 12.95 1 2 .5 0 12.04 22 13.80 13.32 12.84 12.36 11.89 23 13.69 13.23 12.77 ■ 1 2 .3 2 11.86 24 13.64 13.16 12.68 ■ 12,20 11.73 25 1 3 .6 6 13.18 12.70 12.22 11.75 26 13.49 13.03 12.57 12.12 11.66 27 13.59 13.13 12.67 12.22 11.76 28 1 3 .8 0 13.34 12.88 12.43 11.97 29 13.64 13.18 1 2 .7 2 12.27 1 1 .8 1 30 13.78 13.32 12,86 12.41 11.95 31 13.53 13.07 12.61 12.16 11.70

h 4000 5000 5500 6000 6500 7000 n 4500 3 ~~"T4.6t> p n w ™ P'728- ~"I4T1X -r T3V95 ~ ..iy.'62"" 4 13.62 1 3 .3 6 13.10 12.84 12.58 1 2 .3 2 1 2 .0 6 5 13.36 1 3 .0 7 1 2 .7 8 1 2 .5 0 12.21 6 13.07 1 2 .7 8 12.49 12.20 8 1 2 .8 0 9 12-33 257 Table 120. Values of log10ln(h) Computed from $ ^ 2 000km)} and [1^(2000km)}

h n 500 1000 1500 2000 2500 3000 3500 3 "~Y375T~ 15.37 15.21 13". 64"" 14 . W h 15.64 15.38 15.12 14.86 14.60 14.34 14.08 5 15.43 15.14 14.86 14.57 14.28 14.00 13.71 6 15.22 14.93 14.64 14.35 14.06 13.77 13.48 8 14.98 14.68 14.39 14.09 13.79 13.50 13.20 9 14.73 14.40 14.07 13.74 13.41 13.08 12.75 10 14.31 13.97 13.64 13.30 12.96 12.63 12.29 11 14.42 14.05 13.69 13.32 12.95 12.59 12.22 12 14.34 13.96 13.59 13.22 12.85 12.48 12.10 13 14.54 14.13 13.72 13.31 12.90 12.49 12.08 14 14.41 14.00 13.59 13.18 12.77 15 14,31 13.92 13.53 13.14 12.75 16 14.06 13.69 13.31 12.93 12.55 17 14.21 13.80 13.39 12.98 12.57 18 14.22 13.80 13.39 12.97 12.55 19 14.45 14.02 13.58 13.15 12.72 20 14.29 13.86 13.42 12.99 12.56 21 14.27 13.81 13.36 12.90 12.44 22 14.29 13.82 13.34 12.86 12.38 23 14.15 13.69 13.24 12.78 12.32 24 14.12 13.65 13.17 12.69 12.21 25 14.05 13.58 13.10 12.62 12.14 26 13.96 13.50 13.05 12.59 12.13 27 13.89 13.43 12.98 12.52 12.06 28 13.80 13.34 12.89 12.43 11.97 29 13.76 13.30 12.85 12.39 11.93 30 13.82 13.36 12.91 12.45 11.99 31 13.63 . 13.17 12.72 12.26 11.80

h 5500 6500 7000 n 4000 4500 5000 6000 3 x r . f2 1 4 3 5 74739 “ T4723 14.07 13.9*0 4 1 3 .8 2 13.56 13.29 13.03 12.77 1 2 .5 1 1 2 .2 5 5 1 3 .4 2 13.14 1 2 .8 5 1 2 .5 6 12.27 6 1 3 .1 8 1 2 .8 9 12.60 12.31 8 1 2 .9 1 9 12.42 258 Table 121. Values of log10In(h) as Computed from [0n(4000km)} and Cln (4000km)}

h n 500 1000 1500 2000 2500 3000 3500 3 15.86 15.63 15.40 15.17 14.94 4 15.98 15.68 15.37 15.07 14.76 14.46 14.15 5 15.85 15.52 15.20 14.88 14.55 14.23 13.90 6 15.47 15.15 14.84 14.52 14.20 13.88 13.57 8 15.16 14.86 14.57 14.27 13.98 13.68 13.39 9 15.08 . 14.75 14.42 14.09 13.76 13.43 13.10

n h . 4ooo 4500 5000 5500 6000 6500 7000

3 14.71 14.48 14.25 14.02 13.79 13.56 13.33 4 13.85 13.55 13.24 12.94 12.63 12.33 12.02 5 13.58 13.26 12.93 12.61 12.28 6 13.25 12.93 12.62 12.30 8 13-09 9 12.77 APPENDIX E

Balmer Line Intensity Data as Originally Published for Solar Flares and Prominences

259 260 Table 122. Balmer Line Intensities for Prominences as Published by Zirin and Tandberg-Hanssen (i960)

n Quiescent Non-Quiescent 4 15000 15000 6 7000 5000 7 1200 1000 8 . 1000 150

Table 123. Balmer Line Intensities for a Prominence as Published by Elliott, Ellison, and Reid (i9 6 0)

n W'(X) I

3 122 105 3 528 200 6 36 45 7 39 52 Table 124. Balmer Line Intensities in- Prominences as Published in Shih-Huei (1961)

1958 1959

n 11/1 12/4 12/9 3/23 3/30 4/29 6/3 6/24 6/28 7 /1 3 62 70 149 149 213 186 183 112 264 227 4 12 28 59 52 35 31 37 7-4 42 43 5 3-7 11 27 14 16 19 20 2.0 23 24 6 1.9 10 ■6.8 9.6 11 7.5 17 1.0 7-2 13 7 1.3 6.9 4.4 6.6 8.4 5-7 16 0.80 5-6 9-6 8 0.72 4.1 .2 .8 2.5 4.0 2.3 5-4 0.15 2.2 4.1 9 3-3 1.6 1.4 1.3 1.1 4.4 1.1 2-3 10 0 .9 8 0.93 4.7 0.81 1.7 11 1.2 1.0 3.0 . . . .84 1.6 12 - 1.1 1.4 1-9 : .55 .90 13 0.65 O .65 2.3 .55 .84 14 0.37 0.45 T95 262 Table 125. Logarithms of Intensities of Balmer Lines in Prominences as Published by Unsold (1947) and by Athay and Orrall (1957)

Athay and UnsSld (194?) Orrall (1957) log In log In

3 1 5 .5 5 4 2.31 ■ 14.66 5 1.99 14.32 ■ 6 1,066 , 13.99 7 14.12 • 8 1,72 12.91 9 1.24 12.71 10 ' 0.86 12.68 11 0 .8 2 12.68 12 "-6.88. 12.56 13 0.92 12.48 14 0.91 x 12.22 15 . o,4o 1 2 .0 8 16 0.54 12.05 Table 126. Intensities of the Balmer Lines for Solar Flares Seen Against- the Disk as Published in Smith (1963)

h 22"" " ~ 2 2 " 23 ---2T" §3 ---- 53--- m 40 52 01 02 04 07 5 -25 12 20 42 53 20 ■ n 0 0.5?r ' 0 .5 2 0.82 0.52 0.73 0.66 1.91 1.21 1.12 0 .2 7t 0.18 0.48 0.25 0.46 O .39 1.61 0.95 0.86 0.68 Hr 0.54 0.46 0.46 O .6 7 . O .65 2 .2 5 1.24 1.24 0.21 0.12 0.34 0.12 0.33 0.33 2 .0 8 1.14 1.14 c 0.42 0.36 0.57 0.34 0 .5 6 0 .5 0 2.11 -1.06 1.17 b 0.19 0 .0 8 0.29 0 .0 8 0.29 0.18 1.80 0 .7 6 0.84 0.44 0.36 0.61 0.36 2.05 1.14 0& 0.53 0.53 1.05 o;i5 0.05 0.31 0.05 0.22 . 0 .2 3 1.70 0.71 0.79 CS> <=» r-. a=» •0 - 1.80 0 .9 6 1.12 7 0.04 o.o4 0.25 0.05 0.20 0 .2 5 1.48 0.64 0.80 0.24 0.38 0.24 0.36 2.26 Oq 0 .3 1 0.37 1.05 1.20 0 .0 9 0.02 0.16 0.02 0.14 0.14 1.91 0.70 0.85

^Top entry is central intensity, tSeeond entry is emission above the background. Table 126 (Cont.)« Intensities of the Balmer Lines for Solar Flares Seen Against the Disk as Published in Smith (1963)

— gg- — - h "53 23 — 53— 23 . 23 m 16 17 19 21 22 24 27 n s 35 00 56 25 40 05 55

0.82* 1 .2 6 0.92 0-53 0.86 0.75 0.77 0.80 3 0.54f 1.01 0.62 0.21 0.66 0.46 0.48 0.48 0.81 0.82 0.49 0.68 4 0.73 0.77 0.59 0.67 0.38 0.61 0.4? o.i4 0.45 0 .2 6 0.35 0.35 0.62 0.73 O .65 0 .3 5 0 .6 2 0.41 0.46 0 .5 2 5 0.35 0.45 O .38 . 0 .0 9 0.34 0.13 0.20 0.24 O.'SO 6 0.54 0.84 0 .5 8 0.40 0.54 o.4i 0.48 0.25 0.47 0.27 0.08 0.22 0 .1 0 ■ 0 .1 9 0.19 G .6 0 «3S» os» c=o <=>

*Top entry is central intensity, tSecond entry is emission above the background, 264 265 Table 127. Logarithms of the Intensities of the Balmer Lines for a Solar Flare Seen Against the Disk as Published in Svestka (i9 6 0)

Model 1 Model 2 3 15.03 15.69 4 15.16 15.44 5 1 5 .2 0 15.40 6 > 7 15.59 8 15.66 15.55 BIBLIOGRAPHY

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