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THE ASTRONOMICAL JOURNAL, 122 :3335È3360, 2001 December( 2001. The American Astronomical Society. All rights reserved. Printed in U.S.A.
THE SPECTRAL VARIABILITY OF THE CLASSICAL T TAURI STAR DR TAURI1SILVIA H. P. ALENCAR,2,3,4 CHRISTOPHER M. JOHNS-KRULL,5 AND GIBOR BASRI2
Received 2001 April 27 ; accepted 2001 August 14
ABSTRACTWe present the analysis of 103 spectra, collected over more than a decade, of the classical T Tauri star
DR Tau observed with the Hamilton echelle spectrograph at Lick Observatory. The star exhibits strongemission lines that show substantial variety and variability in their proÐle shapes. The emission linesshow signatures of both outÑow and infall, which vary on multiple timescales. The system shows quasi-periodic variations in line intensity and wavelength, but we are unable to recover a unique period thatdescribes all the data. The Balmer and He I line changes are well correlated and appear to result fromreal variations in the accretion and wind Ñows, as opposed to apparent variations caused by changes inthe veiling continuum Ñux. The Balmer line proÐles are generally strongly peaked in the red (vD 100 kms~1) and do not resemble published theoretical magnetospheric accretion proÐles. We suggest that thesystem is seen nearly pole-on. Coupled with a line emissivity that increases strongly near the stellarsurface, this can explain the strongly asymmetric Balmer line proÐles. The Ca II and He I emission-linecomponents are found to be very symmetric and Gaussian in shape, suggesting production in a turbulent(possibly magnetic) region. An additional sporadic high-velocity outÑow component is seen in theBalmer lines and He I. The main characteristic of the lines is their dramatic variability, which indicates avery dynamic interaction between the star and the disk. This is illustrated in several MPEGs animationsshowing the line proÐle variations of DR Tau, which are available on CD-ROM and the World WideWeb.6Key words : stars : formation È stars : preÈmain-sequence
1. INTRODUCTION
Classical T Tauri stars (CTTSs) are young, roughly solar-mass stars that exhibit emission in a wide range of permit-ted lines and occasionally in a number of forbidden lines,together with excess continuum emission from the radiothrough the ultraviolet. Their typical spectral energy dis-tribution is consistent with the presence of a circumstellardisk, and many observations indicate that material from thedisk is actively accreting onto the central star. In addition toaccretion of material onto the star, a strong wind is clearlyseen in the blueshifted absorption features observed in thelow Balmer lines of most stars. The interaction of this star-disk system is believed to explain most of the emission fea-tures seen in CTTSs (Bertout 1989).
One way to study the accretion process is to observe asingle CTTS repeatedly, with a temporal coverage denseenough to sample several phases within a single rotationperiod. With a typical CTTS rotation period of 8 days,nightly sampling provides good coverage. Such synopticobservations permit the analysis of the infall and outÑowdynamics and look for rotational modulation of these sig-natures. The line proÐles of CTTSs tend to be extremelyvariable, suggesting strong intrinsic variation, as well asrotational modulation. As a result, it is a challenge toexplain the variations displayed in even a single line.However, such studies o†er a unique probe of the geometry
ÈÈÈÈÈÈÈÈÈÈÈÈÈÈÈ1 Based on observations obtained at Lick Observatory.2 Department of Astronomy, University of California at Berkeley,
Berkeley, CA 94720-3411 ; basri=soleil.berkeley.edu.3 Departamento de ICEx, UFMG, C.P. 702, Belo Horizonte,F•� sica,
MG, Brazil, 30123-970.4 Current address : Departamento de Astronomia, Instituto
e GeoÐsico, Universidade de Paulo, Avenida MiguelAstronoü mico Sa8 oStefano 4200, Paulo, SP, Brazil,Sa8 o 04301-904 ;alencar=iagusp.usp.br.
5 Space Sciences Laboratory, University of California at Berkeley,Berkeley, CA 94720 ; cmj=ssl.berkeley.edu.
6 See http ://sprg.ssl.berkeley.edu/Dcmj/html/drtau.html.
and astrophysics of the star-disk system. High spectralresolution studies permit detailed analysis of the proÐlemorphology where di†erent components, which probe dif-ferent regions of line formation, can be clearly distin-guished.
DR Tau is one of the best-studied CTTSs. Its visualbrightness increased by several magnitudes during an out-burst in the early 1970s (Chavarria-K. 1979), and since thenit has remained relatively bright (V D 11È12) althoughextremely variable. It shows strong (1È3 mag in V ) photo-metric variability attributed both to cool and hot spots (e.g.,Bertout et al. 1977 ; Bouvier et al. 1993 ; Kenyon et al. 1994 ;Ultchin, Regev, & Bertout 1997) and high veiling values inits optical spectrum (Basri & Batalha 1990 ; Valenti 1994).This results in the near absence of photospheric absorptionlines since the veiling continuum dominates the photo-spheric Ñux at optical wavelengths. DR Tau shows [O I]j6300 in emission with a double peak (one at zero velocityand the other blueshifted ; Hartigan, Edwards, & Ghandour1995 ; Hirth, Mundt, & Solf 1997) suggesting the presence ofa collimated outÑow. The system exhibits quasi-periodicphotometric and spectroscopic variability, with reportedperiods ranging from 4 to 9 days. DR TauÏs emission linesare quite strong and display a wide variety of variableproÐle shapes (Appenzeller, Reitermann, & Stahl 1988 ;Guenther & Hessman 1993 ; Johns & Basri 1995a ; Hessman& Guenther 1997 ; Smith et al. 1999).
In this paper, we present an analysis of the line proÐle(Balmer, Ca II, and He I) variations of DR Tau. Our goal isto explore the nature of the emission-line region, to giveinsight on the relationship between the di†erent line forma-tion regions for these di†erent lines, and to test the generalpredictions of magnetospheric accretion models.
2. OBSERVATIONS
We present the analysis of a sample of 103 spectra of theCTTS DR Tau listed in Tables 1, 2, 3, and 4. The obser-
3335
TA
BL
E1
LIN
EE
QU
IVA
LEN
TW
IDTH
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EL
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FW
HM
OF
TH
EE
MIS
SIO
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UR
BU
LEN
T,A
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BL
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CO
MP
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OF
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TB
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547
529.
664
CA
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0.57
416
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78.3
6787
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2.00
00.
000
208.
661
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1.35
345
.705
...
...
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2..
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247
529.
812
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9.70
482
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0.00
017
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113.
444
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462
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n17
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047
543.
918
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499.
237
111.
693
61.9
092.
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0.00
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347
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9.44
840
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...
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,071
347
546.
801
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6.15
568
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1.84
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262.
694
1.34
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128.
420
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647
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00.
000
267.
055
4.09
8[
112.
423
43.1
79..
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1989
Jan
22,0
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4754
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AT
107.
698
14.3
0294
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67.9
492.
000
0.00
026
7.08
17.
481
[10
6.68
261
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...
...
...
8..
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,072
647
806.
809
CA
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.106
11.5
7311
5.06
664
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2.00
00.
000
238.
649
[1.
210
[17
5.08
369
.440
...
...
...
9..
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89O
ct07
,120
947
807.
008
CA
T62
.733
10.3
7090
.045
76.5
252.
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1.02
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202.
651
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147
807.
812
CA
T68
.112
12.0
0310
2.92
172
.807
2.00
00.
000
197.
967
[1.
007
[19
8.79
6[
29.1
66..
...
...
.11
...
1989
Oct
08,1
216
4780
8.01
2C
AT
70.7
0712
.588
106.
934
71.5
802.
000
0.00
021
7.54
9[
1.12
9[
201.
539
39.7
21..
...
...
.12
...
1989
Oct
09,0
658
4780
8.78
9C
AT
57.2
5410
.489
110.
470
63.4
232.
000
0.00
022
4.09
2[
1.14
6[
185.
798
59.1
81..
...
...
.13
...
1989
Oct
14,0
915
4781
3.88
73
m55
.167
10.6
2511
1.90
671
.747
2.00
00.
000
241.
014
[1.
508
[10
8.22
215
5.54
6..
...
...
.14
...
1989
Oct
26,0
856
4782
5.87
1C
AT
63.4
3312
.096
110.
956
72.5
962.
000
0.00
019
0.35
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1.77
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225.
289
[56
.942
...
...
...
15..
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847
826.
020
CA
T72
.188
13.4
0711
5.38
662
.094
2.00
00.
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254.
896
[1.
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1.55
464
.483
2.47
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26.5
6546
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16..
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347
826.
738
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T51
.683
9.24
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4.29
959
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2.00
00.
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221.
952
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4.41
384
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2.00
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27.8
9150
.366
17..
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447
828.
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CA
T71
.282
11.4
3411
2.15
762
.960
2.00
00.
000
215.
526
[2.
519
[20
4.63
780
.139
3.93
1[
101.
896
90.0
8918
...
1989
Oct
29,1
149
4782
8.99
2C
AT
70.6
0511
.247
105.
393
66.4
532.
000
0.00
021
0.96
0[
1.38
5[
229.
490
69.4
713.
976
[12
5.06
454
.554
19..
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89O
ct30
,053
147
829.
730
CA
T52
.476
7.68
611
3.34
662
.855
2.07
40.
000
181.
547
3.43
6[
144.
120
32.9
76..
...
...
.20
...
1989
Oct
30,1
232
4783
0.02
3C
AT
59.4
968.
862
113.
233
62.2
981.
907
0.00
020
1.30
04.
455
[14
5.66
136
.339
...
...
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21..
.19
89O
ct31
,052
847
830.
727
CA
T55
.971
7.49
511
7.67
863
.627
2.28
60.
000
185.
921
3.24
2[
142.
221
39.7
54..
...
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1989
Oct
31,1
142
4783
0.98
8C
AT
65.3
009.
921
113.
962
65.0
432.
000
0.00
019
8.54
83.
226
[13
7.42
446
.424
...
...
...
23..
.19
89N
ov01
,053
047
831.
730
CA
T65
.894
9.35
011
2.88
968
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2.00
00.
000
222.
458
1.36
8[
100.
008
86.3
78..
...
...
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...
1989
Nov
01,1
238
4783
2.02
7C
AT
78.8
4512
.056
108.
721
65.3
352.
000
0.00
024
3.41
60.
831
[16
2.67
641
.964
3.01
4[
36.3
1146
.054
25..
.19
89N
ov02
,060
247
832.
750
CA
T60
.848
8.52
111
4.44
957
.916
2.00
00.
000
235.
000
[1.
274
[17
7.40
652
.833
3.04
2[
30.7
4172
.738
26..
.19
89N
ov02
,122
447
833.
016
CA
T68
.524
11.1
2010
7.28
965
.288
2.00
00.
000
223.
638
[1.
475
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5.13
333
.843
2.62
5[
51.7
1850
.000
27..
.19
89D
ec01
,045
347
861.
703
CA
T53
.068
10.1
7811
4.47
558
.346
2.00
00.
000
231.
175
[1.
007
[19
7.01
984
.675
...
...
...
28..
.19
89D
ec01
,112
147
861.
973
CA
T64
.817
13.2
5211
2.02
058
.515
2.00
00.
000
235.
245
[1.
008
[18
8.73
861
.943
...
...
...
29..
.19
89D
ec02
,043
847
862.
691
CA
T53
.951
9.93
211
8.54
458
.377
2.00
00.
000
234.
222
[1.
286
[13
5.12
747
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...
...
...
30..
.19
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ec02
,111
147
862.
965
CA
T61
.558
11.7
3812
1.26
060
.064
2.00
00.
000
230.
679
[1.
237
[13
9.37
933
.929
...
...
...
31..
.19
89D
ec03
,042
047
863.
680
CA
T46
.501
8.26
612
5.54
959
.679
2.00
00.
000
231.
309
[1.
438
[12
3.04
074
.459
...
...
...
32..
.19
89D
ec04
,042
347
864.
684
CA
T42
.392
8.65
212
1.49
053
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2.00
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212.
059
[1.
459
[13
2.18
274
.915
...
...
...
33..
.19
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,102
547
864.
934
CA
T59
.841
12.7
5611
9.46
457
.130
2.00
00.
000
221.
889
[1.
194
[13
6.31
766
.429
...
...
...
34..
.19
89D
ec05
,105
247
865.
953
CA
T59
.536
11.9
0912
2.44
357
.557
2.00
00.
000
241.
695
[1.
792
[10
3.36
343
.662
...
...
...
35..
.19
89D
ec06
,041
447
866.
676
CA
T56
.877
11.1
0511
9.93
262
.472
2.00
00.
000
249.
911
[1.
403
[10
2.63
310
7.70
8..
...
...
.36
...
1989
Dec
08,1
124
4786
8.97
73
m48
.992
7.56
411
5.60
371
.507
2.00
00.
000
253.
979
[1.
632
[14
9.24
290
.752
...
...
...
37..
.19
90O
ct21
,105
648
185.
957
CA
T97
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15.7
3411
1.53
570
.736
2.00
00.
000
263.
855
[3.
877
[16
5.16
756
.893
6.76
9[
97.9
9354
.144
38..
.19
90O
ct22
,093
648
186.
898
CA
T80
.654
10.5
9710
7.63
175
.942
2.00
00.
000
257.
718
[3.
590
[18
8.45
250
.634
...
...
...
39..
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ct24
,105
448
188.
953
CA
T77
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14.5
7096
.846
70.0
182.
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260.
399
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2.80
056
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...
...
...
40..
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,111
148
189.
965
CA
T94
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18.9
9610
2.92
367
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680
61.6
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3336
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Con
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HM
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t2C
ent2
FW
HM
2In
t3C
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FW
HM
3In
t4C
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HM
4
41..
.19
90O
ct27
,110
448
191.
961
CA
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5.98
024
.227
91.5
2765
.000
3.00
00.
000
282.
663
[2.
680
[22
1.34
356
.223
8.44
7[
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3747
.717
42..
.19
90O
ct28
,075
748
192.
832
CA
T98
.876
15.8
5490
.812
74.7
602.
000
0.00
028
9.67
5[
2.31
9[
223.
669
63.7
724.
028
[68
.319
49.3
8543
...
1990
Nov
13,0
605
4820
8.75
4C
AT
148.
338
19.8
3610
2.96
762
.124
3.00
00.
000
297.
600
7.87
8[
68.7
8376
.002
...
...
...
44..
.19
90N
ov16
,053
848
211.
734
CA
T76
.426
11.8
1711
8.25
053
.043
3.50
00.
000
249.
020
[1.
515
[10
1.50
764
.721
...
...
...
45..
.19
90N
ov18
,055
148
213.
742
CA
T78
.899
11.9
1210
5.45
468
.656
3.00
00.
000
257.
108
[2.
299
[12
9.31
762
.781
...
...
...
46..
.19
92Ja
n02
,055
248
623.
746
CA
T73
.659
13.5
8886
.466
75.6
722.
000
0.00
026
5.03
7[
2.51
3[
133.
774
82.9
01..
...
...
.47
...
1992
Jan
12,0
724
4863
3.80
9C
AT
64.5
6912
.791
102.
278
57.3
062.
000
0.00
028
6.04
4[
2.31
3[
163.
519
51.8
08..
...
...
.48
...
1992
Jan
13,0
745
4863
4.82
4C
AT
75.3
1012
.370
97.1
2869
.386
2.00
00.
000
302.
952
[2.
272
[18
2.93
334
.050
...
...
...
49..
.19
92Ja
n14
,054
148
635.
738
3m
56.9
068.
585
97.2
9773
.625
2.00
00.
000
234.
008
[1.
889
[19
0.17
533
.008
...
...
...
50..
.19
92Ja
n15
,070
848
636.
797
3m
58.0
949.
766
105.
841
59.8
142.
000
0.00
022
3.96
3[
1.69
8[
188.
339
42.4
061.
999
[54
.556
49.6
4351
...
1992
Jan
15,0
804
4863
6.83
63
m60
.120
10.0
4310
3.73
459
.675
2.00
00.
000
229.
620
[1.
690
[19
0.56
241
.279
2.23
0[
50.5
0747
.975
52..
.19
92Ja
n16
,084
448
637.
863
3m
51.5
707.
900
99.6
5657
.990
2.00
00.
000
220.
378
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514
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1.52
548
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2.41
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57.1
8547
.455
53..
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92Se
p24
,120
148
890.
000
CA
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.557
7.93
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5.68
761
.517
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00.
000
250.
836
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0.03
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...
...
...
54..
.19
92O
ct20
,113
948
915.
984
CA
T65
.300
10.8
7310
2.00
267
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2.00
00.
000
280.
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[1.
961
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5.37
651
.202
...
...
...
55..
.19
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ct23
,103
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918.
941
CA
T62
.360
10.3
5410
3.77
171
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2.00
00.
000
253.
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[16
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...
...
56..
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ct24
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919.
957
CA
T65
.813
10.4
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4.31
468
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00.
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9.02
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...
...
...
57..
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248
920.
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T55
.826
8.01
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1.18
982
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2.00
00.
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416
[2.
207
[16
5.22
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...
...
...
58..
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92O
ct28
,100
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923.
918
CA
T52
.406
9.26
111
3.88
757
.113
2.00
00.
000
285.
637
[2.
413
[16
6.26
557
.915
...
...
...
59..
.19
92N
ov02
,092
748
928.
895
CA
T37
.219
6.36
411
3.32
653
.915
2.00
00.
000
245.
335
[1.
959
[14
7.79
678
.607
...
...
...
60..
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92N
ov14
,113
948
940.
984
3m
52.4
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62.5
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0.00
026
6.23
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1.95
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179.
081
52.0
35..
...
...
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...
1992
Nov
15,0
826
4894
1.85
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m47
.101
8.12
210
2.60
261
.581
2.00
00.
000
244.
306
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[16
9.13
971
.390
...
...
...
62..
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92N
ov16
,111
948
942.
973
3m
59.9
2010
.962
89.2
6467
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000
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928
[18
1.49
648
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...
...
...
63..
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92N
ov23
,081
048
949.
840
CA
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.917
9.11
410
3.13
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000
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1.53
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92N
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,120
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950.
004
CA
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.906
9.27
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8.10
267
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3.22
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...
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92N
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,103
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255
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000
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1.62
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...
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92N
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,072
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952.
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CA
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1.03
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92.4
8674
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000
276.
087
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7.53
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608
18.6
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1992
Nov
26,1
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2.97
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81.5
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93.4
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000
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2.71
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2.31
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800
18.0
8268
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1992
Nov
28,0
606
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4.75
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59.0
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494
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1992
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000
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143.
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1992
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4895
6.80
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000
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142.
829
75.5
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1992
Nov
30,1
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4895
6.97
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AT
50.8
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574
102.
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412.
000
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026
4.84
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147.
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1992
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1992
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1992
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1992
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3337
TA
BL
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Con
tinue
d
EM
ISSIO
NT
UR
BU
LEN
TB
LU
EC
OM
PO
NEN
TS
JDO
BS.
UT
DA
TE
AN
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([2,
400,
000)
TEL.
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.In
t1C
ent1
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HM
1In
t2C
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HM
2In
t3C
ent3
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t4C
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4
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00.
000
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853
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048
[15
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410
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...
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....
1994
Feb
02,0
630
4938
5.77
0C
AT
54.5
6711
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100.
737
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000
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150.
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....
1994
Feb
03,0
648
4938
6.78
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1994
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....
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Feb
15,0
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1998
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19,0
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6.89
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54.3
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....
1998
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....
1998
Dec
16,0
913
5116
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ec17
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251
164.
848
CA
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4311
2.74
565
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218
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1.48
195
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...
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98D
ec18
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551
165.
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CA
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284.
780
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279
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2.11
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...
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....
1998
Dec
18,0
812
5116
5.84
4C
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57.0
7110
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105.
301
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169.
895
CA
T55
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382
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734
[12
6.37
073
.240
...
...
...
NO
TE.È
Lin
eeq
uiva
lent
wid
ths
are
inan
gstr
oms;
cent
erpo
sition
and
FW
HM
are
inki
lom
eter
spe
rse
cond
.Whe
nth
elin
esco
uld
notbe
relia
bly
deco
mpo
sed,
the
equi
vale
ntw
idth
was
notm
easu
red
(em
pty
cells
).
3338
TABLE 2
LINE EQUIVALENT WIDTHS, RELATIVE INTENSITY, CENTER POSITION, AND FWHM OF THE EMISSION, TURBULENT, AND BLUE COMPONENTS OF Hb
EMISSION TURBULENT BLUE ABSORPTION OR EMISSION
OBS. EQW. Int1 Cent1 FWHM1 Int2 Cent2 FWHM2 Int3 Cent3 FWHM3 Int4 Cent4 FWHM4
1 . . . . . . . 3.925 1.311 84.063 74.126 . . . . . . . . . . . . . . . . . . . . . . . . . . .2 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .3 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .4 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .5 . . . . . . . 3.074a 1.910 107.738 44.811 . . . . . . . . . [0.277 [12.607 [36.520 . . . . . . . . .6 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .7 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .8 . . . . . . . 3.633a 2.819 141.254 48.237 . . . . . . . . . [0.615 43.380 75.720 . . . . . . . . .9 . . . . . . . 5.823a 3.053 127.557 57.187 . . . . . . . . . [0.307 [165.113 125.395 . . . . . . . . .10 . . . . . . 4.033a 2.587 136.699 46.486 . . . . . . . . . [0.202 [73.640 108.578 . . . . . . . . .11 . . . . . . 3.125a 0.259 6.614 6.843 . . . . . . . . . 0.163 78.556 96.428 . . . . . . . . .12 . . . . . . 2.609a 1.796 137.021 49.621 . . . . . . . . . 0.136 [150.000 24.546 [0.500 45.737 56.53113 . . . . . . 7.654a 4.048 140.000 59.821 . . . . . . . . . 0.743 [124.102 27.162 [1.176 38.272 62.99114 . . . . . . 7.845a 4.183 142.933 42.297 . . . . . . . . . 0.540 [7.168 59.707 [0.420 [165.081 39.28915 . . . . . . 6.698a 3.750 140.452 42.069 . . . . . . . . . 0.450 [20.472 44.540 [0.323 [163.464 39.93416 . . . . . . 7.811a 3.177 140.679 48.861 . . . . . . . . . 0.818 [31.379 44.355 . . . . . . . . .17 . . . . . . 5.916a 2.541 141.909 41.323 . . . . . . . . . 0.863 [78.388 46.923 . . . . . . . . .18 . . . . . . 6.296a 2.347 141.808 43.026 . . . . . . . . . 1.307 [81.841 41.175 . . . . . . . . .19 . . . . . . 2.872a 1.359 148.329 42.188 . . . . . . . . . 1.130 [86.589 37.391 [0.402 0.000 71.48120 . . . . . . 6.271a 2.277 149.795 44.986 . . . . . . . . . 2.166 [89.395 39.060 [0.484 0.000 67.59821 . . . . . . 4.417a 1.762 153.734 49.913 . . . . . . . . . 1.074 [83.476 44.089 [0.382 0.000 70.00022 . . . . . . 1.036a 0.922 146.798 36.295 . . . . . . . . . 0.366 [86.167 25.391 [0.249 0.000 70.00023 . . . . . . 3.827a 1.841 149.010 43.792 . . . . . . . . . 0.389 [88.645 34.138 . . . . . . . . .24 . . . . . . 4.424a 2.049 140.638 37.723 . . . . . . . . . 0.550 [76.216 57.709 . . . . . . . . .25 . . . . . . 4.544a 2.258 149.337 39.632 . . . . . . . . . 0.435 [80.789 51.630 . . . . . . . . .26 . . . . . . 4.456a 2.719 147.002 35.249 . . . . . . . . . 0.400 [89.283 34.009 . . . . . . . . .27 . . . . . . 5.755a 3.268 105.000 43.296 . . . . . . . . . . . . . . . . . . . . . . . . . . .28 . . . . . . 4.111a 2.240 107.634 41.687 . . . . . . . . . 0.400 [156.426 19.061 . . . . . . . . .29 . . . . . . 3.941a 1.979 119.607 41.389 . . . . . . . . . 0.558 [177.283 27.991 . . . . . . . . .30 . . . . . . 3.364a 1.423 125.129 40.313 . . . . . . . . . 0.621 [157.399 43.802 . . . . . . . . .31 . . . . . . 5.187a 2.507 125.784 43.712 . . . . . . . . . 1.225 [141.101 40.681 [0.419 0.000 75.70332 . . . . . . 2.872a 1.211 126.614 39.977 . . . . . . . . . 0.607 [163.739 37.328 . . . . . . . . .33 . . . . . . 2.132a 0.999 124.641 33.643 . . . . . . . . . 0.749 [163.416 25.167 . . . . . . . . .34 . . . . . . 4.636a 2.033 128.196 46.086 . . . . . . . . . 0.529 [159.802 39.492 . . . . . . . . .35 . . . . . . 6.265a 2.452 121.205 45.001 . . . . . . . . . 0.734 [115.077 61.466 . . . . . . . . .36 . . . . . . 5.793a 3.268 135.321 67.750 . . . . . . . . . [0.765 [28.682 104.949 . . . . . . . . .37 . . . . . . 13.271a 3.001 102.444 58.047 . . . . . . . . . 2.817 [94.717 54.100 . . . . . . . . .38 . . . . . . 13.584a 4.059 107.837 50.207 . . . . . . . . . 2.184 [109.408 64.398 . . . . . . . . .39 . . . . . . 10.628a 5.667 96.100 46.106 . . . . . . . . . . . . . . . . . . . . . . . . . . .40 . . . . . . 10.327a 5.360 111.445 47.567 . . . . . . . . . . . . . . . . . . . . . . . . . . .41 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .42 . . . . . . 23.772a 6.359 83.797 65.141 . . . . . . . . . 3.419 [64.852 49.914 . . . . . . . . .43 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .44 . . . . . . 12.086a 3.588 115.048 44.752 . . . . . . . . . 2.820 [149.734 [50.386 . . . . . . . . .45 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .46 . . . . . . 13.091a 4.261 88.185 55.798 . . . . . . . . . 2.465 [16.707 33.951 . . . . . . . . .47 . . . . . . 10.331a 4.102 102.950 41.439 . . . . . . . . . 1.432 [100.923 60.148 . . . . . . . . .48 . . . . . . 14.063a 3.884 99.589 64.250 . . . . . . . . . 1.993 [101.695 48.608 . . . . . . . . .49 . . . . . . 12.180 2.413 92.344 73.170 . . . . . . . . . 2.185 [97.776 56.318 . . . . . . . . .50 . . . . . . 9.326 1.902 80.967 75.901 . . . . . . . . . 1.868 [127.655 45.809 . . . . . . . . .51 . . . . . . 7.789 1.557 80.719 69.632 . . . . . . . . . 1.497 [110.147 55.899 . . . . . . . . .52 . . . . . . 6.891 1.165 90.388 60.521 . . . . . . . . . 1.361 [92.568 72.935 . . . . . . . . .53 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .54 . . . . . . 4.840 1.889 96.343 45.016 . . . . . . . . . 0.382 [164.531 88.865 . . . . . . . . .55 . . . . . . 5.943 1.898 99.948 49.903 . . . . . . . . . 1.177 [50.919 43.515 . . . . . . . . .56 . . . . . . 6.017 2.113 114.510 54.583 . . . . . . . . . 0.709 [78.789 45.976 . . . . . . . . .57 . . . . . . 6.373 1.698 100.119 50.211 . . . . . . . . . 1.253 [57.018 57.032 . . . . . . . . .58 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .59 . . . . . . 3.391 1.331 104.945 42.702 . . . . . . . . . 0.673 [81.222 39.985 . . . . . . . . .60 . . . . . . 5.204 1.964 102.691 49.069 . . . . . . . . . 0.691 [108.563 45.752 . . . . . . . . .61 . . . . . . 3.600 1.534 95.635 47.957 . . . . . . . . . 0.399 [102.328 37.686 . . . . . . . . .62 . . . . . . 3.976 2.356 95.065 41.445 . . . . . . . . . . . . . . . . . . . . . . . . . . .63 . . . . . . 5.054 1.375 83.661 58.109 . . . . . . . . . 0.803 [99.446 55.639 . . . . . . . . .
3339
3340 ALENCAR, JOHNS-KRULL, & BASRI
TABLE 2ÈContinued
EMISSION TURBULENT BLUE ABSORPTION OR EMISSION
OBS. EQW. Int1 Cent1 FWHM1 Int2 Cent2 FWHM2 Int3 Cent3 FWHM3 Int4 Cent4 FWHM4
64 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .65 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .66 . . . . . . . 3.978 1.538 103.018 43.825 . . . . . . . . . 0.829 [142.665 36.694 . . . . . . . . .67 . . . . . . . 3.278 1.414 108.893 32.168 . . . . . . . . . 0.893 [137.111 39.581 . . . . . . . . .68 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .69 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .70 . . . . . . . 2.535 1.300 97.362 48.137 . . . . . . . . . . . . . . . . . . . . . . . . . . .71 . . . . . . . 4.465 1.932 92.333 56.930 . . . . . . . . . . . . . . . . . . . . . . . . . . .72 . . . . . . . 4.019 1.321 89.856 74.534 . . . . . . . . . . . . . . . . . . . . . . . . . . .73 . . . . . . . 4.431 1.831 96.980 47.699 . . . . . . . . . 0.343 [143.132 64.100 . . . . . . . . .74 . . . . . . . 4.648 1.429 103.121 49.870 . . . . . . . . . 0.509 [136.840 84.082 . . . . . . . . .75 . . . . . . . 4.511 2.038 118.628 48.700 . . . . . . . . . 0.480 [134.634 25.350 . . . . . . . . .76 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .77 . . . . . . . 5.112 1.216 92.971 70.705 . . . . . . . . . 0.645 [68.129 61.069 . . . . . . . . .78 . . . . . . . 5.088 1.312 84.904 69.425 . . . . . . . . . 0.598 [92.647 57.374 . . . . . . . . .79 . . . . . . . 7.414 1.961 97.744 60.457 . . . . . . . . . 1.096 [100.489 58.169 . . . . . . . . .80 . . . . . . . 3.947 1.354 107.173 53.250 . . . . . . . . . 0.306 [108.537 81.592 . . . . . . . . .81 . . . . . . . 1.639 1.593 113.225 47.215 . . . . . . . . . [0.383 [18.600 91.105 . . . . . . . . .82 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .83 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .84 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .85 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .86 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .87 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .88 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .89 . . . . . . . 7.206 3.971 98.661 47.930 0.300 0.000 286.213 [0.867 [94.334 112.982 . . . . . . . . .90 . . . . . . . 3.374 2.242 105.623 48.385 . . . . . . . . . [0.266 [20.865 96.227 . . . . . . . . .91 . . . . . . . 1.754 0.985 113.814 43.854 . . . . . . . . . . . . . . . . . . . . . . . . . . .92 . . . . . . . 6.303 2.820 126.322 54.932 0.400 0.000 281.341 [1.084 [18.825 102.450 . . . . . . . . .93 . . . . . . . 5.599 2.409 122.226 54.488 0.300 0.000 285.356 [0.731 [33.524 105.006 . . . . . . . . .94 . . . . . . . 7.979 3.537 112.018 55.428 0.400 0.000 272.449 [1.098 [39.687 97.791 . . . . . . . . .95 . . . . . . . 8.659 3.153 111.491 61.719 0.400 0.000 301.903 [1.107 [46.815 90.847 . . . . . . . . .96 . . . . . . . 5.565 3.329 107.262 46.102 0.300 0.000 281.164 [0.882 [76.154 111.276 . . . . . . . . .97 . . . . . . . 5.453 3.057 108.984 49.165 0.300 0.000 275.789 [0.882 [76.154 111.276 . . . . . . . . .98 . . . . . . . 6.469 2.979 103.660 48.010 0.400 0.000 275.890 [0.846 [102.025 109.969 . . . . . . . . .99 . . . . . . . 6.764 2.999 105.354 43.873 0.400 0.000 284.762 [0.759 [127.602 102.783 . . . . . . . . .100 . . . . . . 6.466 3.580 100.413 45.944 0.300 0.000 265.130 [0.796 [84.545 123.021 0.300 [119.114 44.741101 . . . . . . 4.166 3.277 101.464 46.216 . . . . . . . . . [0.512 [95.639 120.479 0.300 [127.793 41.007102 . . . . . . 6.271 2.835 112.383 41.565 0.300 0.000 256.836 [0.577 [1.062 85.277 0.300 [188.303 31.668103 . . . . . . 2.975 2.324 115.149 46.113 . . . . . . . . . [0.351 [14.839 96.790 . . . . . . . . .
NOTE.ÈLine equivalent widths are in angstroms ; center position and FWHM are in kilometers per second.When the lines could not be reliablydecomposed, the equivalent width and veiling were not measured (empty cells), and when the Hb line was too close to the CCD edge it was normally cut atabout [250 km s~1, indicated by ““ a.ÏÏ
vations, obtained over the time span of more than a decade,were carried out at Lick Observatory, some using the 3 mShane reÑector, most using the 0.6 m Auxiliary Tele-Coude�scope (CAT) to feed the Hamilton Echelle Spectrograph(Vogt 1987) coupled either to a TI 800] 800 CCD or aFORD 2048] 2048 CCD. The entire spectral format is notcovered with the smaller CCD, so observations were gener-ally obtained in one of two settings : (1) a red setting cover-ing 52 partial orders from D4900 to D8900 and (2) a blueA�setting covering 38 partial orders from D3900 to D5200 A� .Whenever possible, blue and red observations wereobtained in the same night or on successive nights. Thelarger CCD, installed in 1992, records D92 orders coveringthe optical spectrum from D3900 to D8900 The meanA� .resolution of the spectra is j/*j B 48,000, and the exposuretimes varied from 45 minutes to 1 hour 15 minutes, depend-ing on the telescope and the CCD used.
The reduction was performed in a standard way
described by Valenti (1994), which includes Ñat-Ðelding withan incandescent lamp exposure, background subtraction,and cosmic-ray removal. Wavelength calibration is made byobserving a ThAr comparison lamp and performing a two-dimensional solution to the position of the thorium lines asa function of order and column number. Radial and bary-centric velocity corrections have been applied, and all thedata shown here are in the stellar rest frame. No Ñux cali-bration has been attempted for these spectra ; rather eachspectrum has been continuum normalized. Because of dif-ferences in weather conditions, exposure times, telescopeused, and efficiency between the di†erent CCDs, there is awide range in the signal-to-noise ratio (S/N) in the data.
3. MEASUREMENTS
3.1. Survey of L ine ProÐlesA sample of the various line proÐles discussed in the
TA
BL
E3
LIN
EE
QU
IVA
LEN
TW
IDTH
S,R
ELA
TIV
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SPECTRAL VARIABILITY OF DR TAURI 3347
following sections is shown in Figure 1. The Ðrst two obser-vations are separated by 3 days in October 1989 and the lasttwo by 2 days in December 1998. The Ha line presents typeIVB P Cygni proÐles (Reipurth, Pedrosa, & Lago 1996) inD75% of the observed DR Tau spectra (e.g., Fig. 1cÈ1d).Sometimes, however, a blue shoulder develops (e.g., Fig. 1b),and the Ha line resembles the proÐles of stars such as RWAur : two very symmetric, well separated peaks (Johns &Basri 1995a ; Muzerolle, Hartmann, & Calvet 1998b ;Alencar & Basri 2000). An intermediate situation is shownin Figure 1a where a blue bump has just appeared and isstill very close in velocity to the red emission. The Hb linemimics the behavior of Ha, with smaller overall emissionintensity, particularly in the red emission peak.
As described below, the He I line can generally be decom-posed into broad and narrow Gaussian components. TheHe I line characteristics vary substantially, as is illustratedin Figure 1, where the narrow component (NC) increases instrength from Figure 1c to Figure 1d over 2 days while thebroad component (BC) becomes signiÐcantly narrower atthe same time. The He I proÐle appears to change in acorrelated manner with the Balmer lines. We notice a blue““ bump ÏÏ in the He I BC when the blue emission in theBalmer lines is the strongest (Fig. 1b).
The variation in the Ca II lines is primarily one of inten-sity, and the line does not show any signiÐcant change inshape when the Balmer proÐles develop blue emission
peaks. The BCs of the Ca II lines are much narrower at theFWHM (*vD 42 km s~1) than the He I BCs (*vD 90 kms~1), possibly indicating that the He I BC originates in amore turbulent region, such as the accretion hot spot, or ina region with higher velocity gradients than the Ca II BC.
In Figures 2 and 3 a sequence of spectra obtained inOctober 1989 is shown. In this sequence, rapid variation inthe Ha proÐle is observed, including the development of ablue bump, which then fades away, all in less than a week.In addition, a new blue emission shoulder just begins toappear in the last two proÐles. The Hb proÐle shows thesame type of variation, with the blue emission becoming asstrong as the red emission peak. During this time the He I
line displays small changes in strength and shape, and aredshifted absorption component is just visible in some ofthe spectra, while the Ca II line only appears to change inintensity. Another week-long series of Ha proÐles fromDecember 1998 is shown in Figure 4. In contrast to the lastsequence, this time series shows almost no variations in theproÐle shape, which remains P Cygni as in the entire run.
MPEG animations showing the variability of each line,as well as animations showing the simultaneous variationsof many of the lines, can be seen on the World Wide Web7and are also available on CD-ROM.
ÈÈÈÈÈÈÈÈÈÈÈÈÈÈÈ7 See http ://sprg.ssl.berkeley.edu/Dcmj/html/drtau.html and are also
available on CD-ROM.
FIG. 1.ÈSample of the observed line proÐles. The horizontal dotted lines represent the continuum level at 1. The vertical dotted lines trace the stellar restframe at 0 km s~1. The numbers in the leftmost plots give the Julian Dates of the observations minus 2,400,000.0.
3348 ALENCAR, JOHNS-KRULL, & BASRI Vol. 122
FIG. 2.ÈHa (left) and Hb (right) sequence of observations from October 1989. JD is the Julian Date minus 2,400,000.0, and the dotted lines have the samemeaning as in Fig. 1.
3.2. ProÐle DecompositionMost of the lines appear to be composed of more than
one distinct component, which displays di†erent character-istics and variability. Therefore, we decomposed the lineproÐles into Gaussian component parts to investigate theirbehavior separately and to be able to better compare theo-retical magnetospheric accretion line proÐles with the com-ponents of the line supposedly formed in this region. Wechoose Gaussian components because, as shown below,they provide a good description of the line proÐle with aminimum number of separate components.
We decomposed the Ha P Cygni proÐles primarily into(1) a strong redshifted emission peak, presumably due to theaccretion Ñow, (2) a blueshifted wind absorption com-ponent, and (3) a relatively low amplitude component withlarge FWHM centered at rest velocity (Fig. 5, left). We favorthis decomposition instead of a (more than twice as strong)emission component (likely due to accretion) centered atrest velocity plus a (sometime multiple) strong blueshiftedwind component (Fig. 5, right) for several reasons. Thedecomposition with a centered emission component thatÐts the observations always required very strong windlikeabsorption components, with one of them slightly red-shifted. It also appears that the upper red wing of theobserved emission is rather steep and is not as well Ðtted bya broad Gaussian emission component centered at rest as itis by a redshifted narrower emission component. The peakof the observed redshifted emission is very symmetric, and
this is hard to explain if each wing were inÑuenced by di†er-ent processes : accretion for the red side and a wind for theblue side. We also obtain a better Ðt with our adopteddecomposition (s2\ 3.5 ; Fig. 5, left) than we do decompo-sing the proÐles with a centered emission component(s2\ 12.3 ; Fig. 5, right), using the same number of freeparameters. The s2 values reported are reduced s2. Finally,in ° 4.3, we show that the variability analysis also supportsthe adopted decomposition, and our strongest reasons forthis preference are discussed in that section. These includethe fact that the redshifted emission component behavestemporally like a single coherent line (similar to He I), butincoherently with respect to the blue wing, so that we Ðnd itunlikely they are related by a single emission component.
Muzerolle, Hartmann, & Calvet (2001) obtained a goodÐt to the NaD lines of DR Tau by using their new magneto-spheric accretion models and an almost edge-on conÐgu-ration (i \ 70¡). Using the same stellar and magnetosphericparameters they cannot Ðt the observed Ha proÐles. Thesemodels produce a small, rather centered emission line forHa from the funnel Ñow and they suggest that most of theHa emission comes from a hot wind. In that case the correctline decomposition would be the one shown in Figure 5(right), where the wind would be responsible for both emis-sion and absorptions.
According to our proposed decomposition, when a blueshoulder appears in the Balmer lines, an extra blue emissioncomponent is necessary to Ðt the proÐles (15% of the cases
No. 6, 2001 SPECTRAL VARIABILITY OF DR TAURI 3349
FIG. 3.ÈHe I (left) and Ca II (right) sequence of observations in October 1989. JD is the Julian Date minus 2,400,000.0, and the dotted lines have the samemeaning as in Fig. 1.
as in Figs. 1a and 6, left) and the blue absorption may evendisappear completely (11% of the cases, as in Figures 1b and6, right). Ardila et al. (2001) also show evidence of extra bluecomponents in the far-ultraviolet lines.
The Hb proÐles resemble the Ha ones, except that thebroad and low Gaussian contribution does not seem to bevery important for this line ; however, our attempts tomeasure it are hampered by the lower S/N available at Hb.
The He I lines clearly display a broad component and anarrow component (see also Beristain, Edwards, & Kwan2001), and in D20% of our spectra a redshifted absorptioncomponent at D250 km s~1 is also visible (see Fig. 1c).Both broad and narrow components are generally verysymmetric and well Ðtted by Gaussians. The NC is nor-mally found centered at the stellar rest velocity while the BCis generally slightly blueshifted (again, see also Beristain etal. 2001).
The Ca II infrared triplet lines show basically only abroad component, although the presence of a small NC canbe inferred because of the somewhat triangular shape of theproÐles. In this particular case it is very hard to disentangleeach component, and we Ðt the Ca II line with only a singleGaussian in emission. As noted before, the BC FWHM inCa II is about half that of the He I line. All the Ca II infraredtriplet lines have a contribution from a Paschen emissionline in the high-velocity region of their red wings that iseasily identiÐed.
All lines were decomposed as described above. The equiv-alent width measurements and the Ðt parameters are givenin Tables 1, 2, 3, and 4.
3.3. VeilingDR TauÏs photospheric spectrum is highly veiled, as has
already been noted by several authors. Published veilingvalues range from 0.7 to 20 (Basri & Batalha 1990 ; Valenti,Basri, & Johns 1993 ; Guenther & Hessman 1993 ; Beristainet al. 2001). Ideally, to obtain a good veiling measurementrequires at least a couple of photospheric lines clearly dis-tinct from the continuum. Unfortunately, at our S/N, theonly strong photospheric signature that is present in amajority of the spectra is the Li I j6707 line, which is farfrom being ideal for use as a probe of the veiling. Stout-Batalha, Batalha, & Basri (2000) showed that in the case ofRW Aur, the Li I deveiled line strength is enhanced byaccretion by more than a factor of 2 between the lowest(V \ 0.3^ 0.3) and highest (V \ 6.1^ 1.7) veiling values.We conÐrm a similar trend in DR Tau by using only ourhighest S/N data where the veiling can be determined fromphotospheric lines other than Li I. Nevertheless, since theLi I line is our only option in most cases, we use it to measureveiling values (ranging between 3 and 6) in the spectrawhere they could be measured reliably (the ones with higherS/N) and checked for consistency with other very faintphotospheric lines where available. In principle, we may be
3350 ALENCAR, JOHNS-KRULL, & BASRI Vol. 122
FIG. 4.ÈHa sequence of observations in December 1998. JD is theJulian Date minus 2,400,000.0, and the dotted lines have the same meaningas in Fig. 1.
missing the high veiling values, as they could be unde-tectable in the low S/N spectra, but given the Li I corre-lation with veiling a highly veiled line may still be visible.
To calculate the veiling we compared the Li I line of an
M0 V standard star with that of DR Tau. We Ðrst rotation-ally broadened the continuum-normalized spectrum of thestandard star and then applied the veiling to the standardÏsspectrum, comparing it with the observed DR Tau spec-trum until a good match was found. The best adjustmentwas determined by eye. Whenever we could use photo-spheric lines other than Li I a unique veiling value did notperfectly match all the lines and the results presented corre-spond to a mean value obtained with all the photosphericlines available. The veiling values obtained are listed inTable 4. They are used to search for correlations betweenveiling-corrected line equivalent widths, deÐned as inJohns-Krull & Basri (1997) :
W eq0 \ Weq(V ] 1) , (1)
where is the measured equivalent width and V is theWeqveiling. The veiling values were computed between 5000 and7000 A� .
4. ANALYSIS
4.1. ProÐle VariabilityWe show in Figure 7 the average line proÐles for the
strong emission lines of DR Tau, which are obtained bytaking the mean value in each velocity bin of a given lineover all the observations. Also shown in the shaded area arethe normalized variance proÐles, which measure theamount of variability in each velocity bin in the line proÐle.As deÐned by Johns & Basri (1995a), the temporal varianceof the line at each velocity bin is given by
;v
\ ;i/1n (I
v,i [ Iv)2
n [ 1, (2)
where n is the number of observations, is the proÐleIv,iintensity at a given velocity (v) in each observation (i) and I
vis the mean intensity at a given velocity (v) over all theobserved proÐles. The normalized variance proÐle is thevariance proÐle as deÐned above divided by the average lineproÐle.
Although the proÐles in Figure 7 are not corrected for
FIG. 5.ÈHa decomposition. Top, Dashed lines showing the individual components ; bottom, dashed line representing the Ðnal Ðt (adding all the individualcomponents). The dotted line is the red side of the proÐle reÑected about zero velocity. L eft, decomposition scheme generally adopted in this paper ; right, Hadecomposition, assuming a very large emission component centered at rest velocity. The number at upper left is the Julian Date of the observation minus2,400,000.0.
No. 6, 2001 SPECTRAL VARIABILITY OF DR TAURI 3351
FIG. 6.ÈDi†erent Ha proÐles. All the lines have the same meaning as in Fig. 5, but at right the dotted line is the red side reÑected about 20 km s~1. Thenumber at upper left is the Julian Date of the observation minus 2,400,000.0.
veiling, it is apparent that much of the variability in the lineproÐles is not simply due to variations in the veiling itself. Ifthis were the case, the shape of the variance proÐle would bethe same as the original line proÐle itself since the entire linewould change simultaneously. Most of the variability in Hais in the blue side of the proÐle, which we interpret asresulting from changes in the wind. The red emission peakshows little variation relative to the blue side, indicating the
two parts of the proÐle are not primarily inÑuenced by thesame processes. The Hb variability mimics that of Ha,except that the sharp cuto† on the blue is due to the factthat some of the Hb spectra used fall too close to the leftedge of the CCD and are truncated in the blue wing. TheHe I and Ca II lines do not show strong variations in generalbut can vary sporadically. Figure 7 clearly shows the inÑu-ence of the redshifted absorption at D250 km s~1 in He I
FIG. 7.ÈAverage line proÐles (solid lines) and variance proÐles (gray shaded areas)
3352 ALENCAR, JOHNS-KRULL, & BASRI Vol. 122
probably due to the infall of material at nearly free-fallvelocities in the accretion Ñow. The He I line is the weakestline studied and therefore the most inÑuenced by veilingvariations. Its normalized variance proÐle is consequentlythe one that most resembles the average line proÐle.
4.2. Periodogram AnalysisMany attempts have been reported in the literature to
determine the rotation period of DR Tau. Richter et al.(1992) obtained a marginal detection of a 9.0 day periodusing di†erential CCD photometry. Bouvier et al. (1993),analyzing broadband light curves, suggested the presence oftwo periods of 2.8 and 7.3 days. Kenyon et al. (1994) foundphased optical and near-IR photometric variations withtimescales of 5 or 10 days. However, Kenyon et al. (1994)argue that neither theirs nor BouvierÏs sample is longenough in time to demonstrate periodic behavior. Bouvieret al. (1995) observed DR Tau again (as part of theCOYOTES campaign) and were able to conÐrm the 7.3 dayperiod they had found earlier. They also reported a 9.0 dayperiod that agrees with the determination by Richter et al.(1992), but this value has a larger uncertainty as their lightcurve covers less than two photometric periods. Johns &Basri (1995a), analyzing the Ha proÐle variations of DRTau, obtained a period of 5.1 ^ 0.2 days and another pos-sible period of 7.8^ 0.4 days. Hessman & Guenther (1997)analyzed a series of spectroscopic and photometric obser-vations and found that the strength of the strong emissionlines of DR Tau seemed to be quasi-periodic in time. Theyobtained a 4.48 day period, mainly due to the Balmer linevariations and a 4.5È5.0 day period, mostly due to Ca II. Aspointed out in ° 3.3, the veiling of DR Tau is high and quitevariable ; thus, the photometric periods obtained because ofhot spots could just be reÑecting quasi-periodic variationsin the accretion source and not necessarily the stellar rota-tion itself (Johns & Basri 1995a).
We performed a periodogram analysis on the intensityvariations of the di†erent lines using the Scargle (1982)periodogram estimator as modiÐed by Horne & Baliunas(1986), which is appropriate to handle irregularly spaceddata. We chose initially to investigate small groups of datainstead of the entire set at once since there are occasionaltime gaps of several months in our sampling. In addition,because of the fact that several di†erent periods have beenreported for this star, it seems prudent to analyze the dataon a month-to-month basis. In general, results of the perio-dogram analysis yielded periods of 4 to 9 days in many ofthe data blocks for the di†erent lines, but we do not recovera unique period that describes the variations of all the lines.We also tried to use the complete data set, but we could notretrieve any signiÐcant period from it.
We note in Figure 4 that the Ha red emission peak posi-tion seems to vary periodically. However, as in the case ofthe line intensity, although the data may look ordered andperiodic in certain epochs, it is quite random most of thetime. We show in Figure 8 the periodogram results for 1998December and for 1989 October, which correspond toFigures 3 and 4. In these same Ðgures the dotted lines showthe power spectrum of random noise sampled at the sametimes as the data. This indicates that the peaks that appearin the 1989 October data are likely to be due to the sam-pling rather than real periodicities. On the other hand, the1998 December data suggest that a 9.3 day periodicity ispresent in the data.
4.3. Correlation Matrices
To investigate how the proÐle variations are correlatedacross a given emission line we calculated autocorrelationmatrices for each of the main emission lines of DR Tau.Correlation matrices are shown for each line in Figure 9,where the lowest plotted contours correspond to 99.9%conÐdence level for the presence of a correlation. We nextperformed tests to explore the inÑuence of veiling variationson the correlation matrix. For example, we added to theaverage Ha proÐle a random veiling value in the range3 \ V \ 5, which corresponds to the typical values weobtained for DR Tau, and we computed the correspondingcorrelation matrix for this simulated data. This produced alarge square region of positive correlation centered in thematrix at [200 km s~1\ v\ 400 km s~1. The corre-lation matrix computed from the observed Ha proÐles, onthe other hand, shows substantial structure along the diago-nal that cannot be due simply to variability in the veiling.We note that the region corresponding to the wind com-ponent ([300 km s~1\ v\ [100 km s~1) does not cor-relate with any other part of the proÐle. This stronglysuggests that the blueshifted part of the line is produced in asigniÐcantly di†erent region than the emission peak. Theblue side comes from a region that varies substantially(according to the variance proÐles) and not uniformly, sincethe correlation matrix is very narrow in this part of theproÐle. The Hb matrix also does not show a squarish struc-ture. The correlation in this case is only along the diagonal,indicating little relation between the variations along theline proÐle.
The Ca II line exhibits a simple proÐle with just one com-ponent and is probably formed in a small region thatresponds quickly to any changes in emission properties. Thecorrelation matrix of the Ca II line is by far the most squar-ish we obtained, consistent with this interpretation, but alsoconsistent with variability due to veiling changes. However,if veiling variations were responsible for the Ca II variabil-ity, it should also appear in the matrices of the other lines.The He I matrix shows that the emission parts of the line arewell correlated, either because of veiling or because of theaccretion process itself, while the redshifted part of theproÐle does not show much correlation with the emissioncomponents. Again, if changes in the veiling were primarilyresponsible for the apparent emission-line variability, theentire line would show a good correlation, not only theemission part of it. From the correlation matrices we con-clude that although the veiling is important, intrinsic lineproÐle variability is primarily responsible for the observedvariations.
We note that the red and blue wings of the Ha emissionpeak are well correlated, indicating that the same process isa†ecting both sides simultaneously. Since the core of thewind component itself is not correlated (or anticorrelated)with any other part of the line, we Ðnd it unlikely that thewind is responsible for the shape of the blue side of theemission peak at the same time the accretion Ñow isresponsible for shaping the red side. In other words, webelieve these correlations do not support decomposing theHa proÐle into a large emission component centered at restwith several absorption components representing the wind.
Another interesting use of the correlation matrices is tosearch for correlated variations among di†erent lines, whichmay point to a common region of line formation. We calcu-
No. 6, 2001 SPECTRAL VARIABILITY OF DR TAURI 3353
FIG. 8.ÈPower spectrum of the Ha red emission peak center in 1989 October and 1989 December. The solid lines correspond to the observed data, andthe dotted lines, to random noise sampled at the same times as the data.
lated the cross-correlation matrices between the primaryemission lines and found that Ha and Ca II do not appear toshow any correlation at all. The emission part of the He I
line shows a good correlation with most of the Ha emissionpeak, implying that they are both formed at least partiallyin the same region, likely the accretion Ñow itself, accordingto the magnetospheric models. The He I BC is correlatedwith the Ca II emission (which is basically composed only ofa BC). They are both thought to be produced in the funnelÑow, although in somewhat di†erent regions. The He I lineis formed in high-temperature (T [ 20,000 K) regions, whilethe Ca II is thought to come from low temperatures (4000K \ T \ 7000 K). The correlation shows that, despite thedi†erent temperatures characteristics of their formationregions, there is some common process a†ecting them both.
The Ha and Hb lines present a strong correlationbetween each other, as expected, since they are supposed tobe formed basically in the same regions. The correlation ismainly along the diagonal, showing that the same com-ponents of the two lines are well correlated with each other.The He I line correlates with the Hb emission peak, onceagain showing that the Balmer lines and the He I must sharesome common formation region. It appears from the matrixthat the He I NC tends to correlate better with the central
part of the Hb emission peak, while the BC shows a strongcorrelation with the peakÏs red wing. Di†ering from Ha, theHb red emission peak correlates with the Ca II emission line.
We also checked for correlations with 1, 2, and 3 day timelags between the di†erent lines. We found that the He I linecorrelates well with the Ha emission peak for a 1 day lag(Ha leading He I). Ha probes a much larger volume thanHe I, which may be produced, at least partially, in the shockregion. Both lines also show a strong correlation with notime lag, while the correlation with a 1 day lag suggests thatvariations in the accretion Ñow Ðrst appear in Ha, becauseof the larger volume probed by Ha compared with that ofHe I. No correlation was found when Ha lags He I by 1 day.The Ca II line slightly correlates with Hb when Ca II lags by1 day. They also correlate with each other with no time lag.Nothing was found when the Ca II line leads Hb by a day.The Ha and Hb lines slightly correlate with each otheralong the main diagonal with a lag of 1 day for either lineleading. The variation of both lines in time is very similar,and they would be expected to show some time-laggedcorrelation with each other. We could not Ðnd any corre-lation with a 2 or 3 day time lag. The number of data pointssuited for this study is actually quite high (36 data pairs inthe Ðrst case and 32 pairs in the second) ; therefore, we
3354 ALENCAR, JOHNS-KRULL, & BASRI Vol. 122
FIG. 9.ÈAutocorrelation matrices
believe the failure to see correlations with these lags indi-cates there really is no such correlation as opposed to a lackof sufficient data to demonstrate it.
5. DISCUSSION
5.1. Comparison with Other CT T SsThere are only a few synoptic studies of CTTSs in the
literature. Although most results tend to favor the magneto-spheric accretion model predictions, some also point todeÐciencies present in this model. SU Aur and Sz 68 are twoCTTSs that strongly support the magnetospheric accretionmodels. Johns & Basri (1995b) analyzed the emission-linevariability of SU Aur and showed that Hb presented simul-taneous periodic variations in both infall and outÑow signa-tures, which were 180¡ out of phase with each other. Johns& Basri (1995b) suggested this could be explained by themagnetospheric accretion and wind model of Shu et al.(1994) if the dipole axis of the magnetosphere is inclinedwith respect to the stellar rotation axis. Johns-Krull &Hatzes (1997) studied the CTTS Sz 68 and also found evi-dence for rotational modulation of the accretion Ñow astraced by the Ha proÐle. The Doppler image produced byJohns-Krull & Hatzes (1997) revealed a slightly o†-center
polar spot that could also be due to an inclined dipolar Ðeldcomponent. However, analysis of emission-line variabilityin other stars (Johns & Basri 1995a ; Johns-Krull & Basri1997) does not generally show periodic variations as mightbe expected from the magnetospheric accretion model.
The line proÐle variations of DF Tau, the star that orig-inally motivated the idea of magnetospheric accretion(Bertout, Basri, & Bouvier 1988), was found to only par-tially conÐrm the general predictions of the magnetosphericaccretion model (Johns-Krull & Basri 1997). Periodicmodulation in the line proÐles appears only intermittently,and the periodic variations do not appear in the expectedregions of the proÐle. Johns-Krull & Basri (1997) alsosuggest an important role for large-scale turbulence in theproduction of the broad symmetric Balmer line wings ;however, Muzerolle et al. (2001) suggest that Starkbroadening can explain the observed wings if the density inthe accretion Ñow is high enough. Recently Ardila et al.(2001) have called this into question, since the Mg II linesshow essentially the same broadening as Ha.
Gullbring et al. (1996) analyzed the line proÐle variationsof the CTTS BP Tau. They found that most emission linesshow variations in both equivalent width and proÐle shape
No. 6, 2001 SPECTRAL VARIABILITY OF DR TAURI 3355
on timescales from a few hours up to a few days. They foundthat a signiÐcant part of the line emission was consistentwith a magnetospheric accretion model with an inclineddipole Ðeld. However, they also show that the Balmer linestend to be substantially more symmetric than predicted inthe theoretical proÐles available at the time for magneto-spheric accretion models. They propose the presence of astrong wind component to help explain the symmetry of theextended line wings.
Unlike some of the stars discussed above, DR Tau doesnot show a clear periodicity in our analysis. This lack ofperiodicity could either be due to its dipole axis beingaligned with the rotation axis or simply to a low inclinationof the system, which would require a very inclined dipoleaxis to produce detectable periodicities as the star rotates.
We notice that DR Tau presents a relatively intense Haline variability when comparing the normalized varianceproÐles of Figure 7 with the results of other stars in Johns &Basri (1995a) and Johns-Krull & Basri (1997). Other emis-sion lines of DR Tau, however, show variability rathersimilar to stars such as DF Tau. In general, the level ofvariability of DR Tau is not very di†erent from otherCTTSs, although DR Tau presents a dramatic diversity ofBalmer line proÐles.
Most CTTSs in Johns & Basri (1995a) present Ha corre-lation matrices that are mostly diagonal, indicating that thechanges within the line are not correlated. The Ha corre-lation matrix of DR Tau shows a good correlation betweenthe di†erent parts of the line except for the wind region anddoes not resemble the matrices of most stars in Johns &Basri (1995a). If the DR Tau system is seen at low inclina-tions, the Ha emitting area should not change much as thestar rotates, and we would therefore expect (as observed) agood correlation between the components of the proÐle thathave a common origin, like the redshifted emission and thefar emission wings that both come from the accretion Ñowin our decomposition scenario. Low system inclination isalso the case for T Tau, shown in Johns & Basri (1995a),which displays an Ha matrix very similar to DR Tau. TheHb line of DR Tau, on the other hand, produces a corre-lation matrix that is very diagonal, indicating little relationbetween the variations in the di†erent positions in the lineproÐle. This line was basically decomposed into one emis-sion component due to accretion and one absorption com-ponent due to a wind, which we would expect to show somecorrelation with each other, following the Shu et al. (1994)model. However, the results obtained with the correlationmatrices of Ha and Hb do not show the expected corre-lation between the wind and accretion components.
The most common type of Ha proÐle presented by DRTau (type IVB) corresponds to less than 5% of the observedline proÐles of T Tauri stars in the Reipurth et al. (1996)atlas, and the Ha line possesses wings that extend as far as^500 km s~1, also a rarity in this atlas. Both characteristicsshow that the Ha proÐle of DR Tau is not common amongCTTSs. The He I line proÐles with the BC generally blue-shifted, on the other hand, are typical of highly veiledCTTSs, according to Beristain et al. (2001).
5.2. System GeometryThe geometry of the DR Tau system, which produces
such variable and diverse line proÐles, is puzzling. It is verydifficult to estimate even the inclination of the system, sincewe do not know precisely the rotation period of the star.
Kenyon et al. (1994), using a period of 5 days andv sin i ¹ 10 km s~1 and assuming estimatedR
*\ 2 R
_,
that the system inclination should be i ¹ 30¡ (they actuallyobtained i ¹ 60¡, but there was an error in theircalculations). This also agrees with the spectral energy dis-tribution Ðts they determined for DR Tau: the system ismore likely viewed closer to pole-on than edge-on. Theyfurther estimated that i ] b B 95¡, which would imply thatb, the angle between the magnetic and rotation axes, isº65¡. However, their values for i and b depend on uncer-tain parameters, most notably the stellar radius and rota-tion period, as well as the illumination pattern of theaccretion ring, which is likely to be nonuniform for tiltedmagnetospheres (e.g., Mahdavi & Kenyon 1998).
In agreement with Kenyon et al. (1994), we suggest thatthe shapes and variability of the permitted emission lines inDR Tau are most easily explained in a more pole-on thanedge-on geometry. However, unlike Kenyon et al. (1994), werequire that the dipole component of the magnetic Ðeld isaligned closely with the stellar rotation axis. We assumethat most of the time the high-velocity portion of the mag-netospheric Ñow (closest to the star) is along poleward-oriented Ðeld lines and is therefore directed away from theobserver (redshifted). This portion of the Ñow is generallyvisible, giving rise to the strong redshifted emission peak,while the corresponding blueshifted portion of the Ñow isapparently occulted by the star, resulting in a weak orabsent blueshifted emission peak (Fig. 10). Part of the low-velocity redshifted Ñow is also occulted by the disk in thisconÐguration, which enhances the Balmer lineÈredshiftedasymmetry. In this picture, we require a strong concentra-tion of the visible brightness of the funnel Ñow to thoseregions very close to the stellar surface. The Ha decomposi-tion we adopted based on this picture provides a better Ðtto the observed line proÐles than one with a very largecentered emission component, which would represent asituation where the Balmer line forms throughout thefunnel Ñow with no particular enhancement near the star.
5.3. Ha Emission : W ind Origin?As mentioned in ° 3.2, Muzerolle et al. (2001) suggested a
very di†erent scenario to describe the DR Tau system. Theysuggest that DR Tau is seen at high inclination (i\ 70¡) andmost of the Ha emission comes from a hot wind, only asmall part being due to accretion. Their new accretionmodels Ðt well the observed NaD lines of DR Tau butproduce small, rather centered Ha emission with the sameparameters.
The Ha proÐle does have some resemblance to what isexpected from a spherical wind. This was the original model
FIG. 10.ÈSimpliÐed representation of the DR Tau system, showing theÑow of material from the disk (arrows).
3356 ALENCAR, JOHNS-KRULL, & BASRI Vol. 122
for the Balmer lines in T Tauri stars (TTS), and modelproÐles are discussed by Hartmann et al. (1990), forexample. These models were discarded because they reallydo not work for most TTSs (e.g., Johns-Krull & Basri 1997),but DR Tau has a rather atypical Balmer proÐle, whichclearly shows wind absorption. The redward asymmetry ofthe emission results in a normal P Cygni scenario, and someof the proÐles in Hartmann et al. (1990) bear more than apassing resemblance to our observations. Nonetheless wedo not favor this as the best explanation for them for anumber of reasons.
To get the right proÐle width, Hartmann et al. (1990) hadto posit a very large turbulent component to the wind.While we would agree that turbulence is probably thesource of much of the broadening, it is much easier toimagine a source for such strong broadening in the stellarmagnetosphere than in an extended wind. In the wind inter-pretation, the blue edge of the red emission peak is due towind absorption, so its behavior should be correlated to thedepth and velocity of the main blue absorption feature, andnot with the red wing of the emission peak. That is not whatwe see ; the red emission peak tends to maintain symmetrybetween its red and blue side, in the face of dramaticchanges in the wind absorption. We also see instances inwhich there appear to be two separate wind components(one at low and one at high velocity), and the red emissionpeak does not correlate with either of them. The blue emis-sion peak in the wind models comes from partialuncovering of a broad central emission proÐle, and that isclearly not the correct explanation for the occasional bluepeak in our observations, whose strength is basically uncor-related with either the velocity or the strength of the windabsorption. In this case, one might get around the issue byaccepting our explanation that this occasional blue peak isdue to high-velocity outbursts. In the wind model, however,the emission would have to occur exterior to the site of windabsorption.
To better distinguish a generally spherical wind from thehot accretion column plus collimated wind that we favor,one would have to try harder to Ðt the data with sphericalwind models. Such models must be able to Ðt the variety ofobserved Ha proÐles, as well as the other strong lines (likethe Ca II and Mg II lines). Furthermore, the wind wouldhave to be able to change dramatically on a timescale of onenight or less, as seen in the proÐles. At this juncture, we leantoward the idea that the red Ha emission is produced ratherclose to the star and that its redward predominance is dueto inÑow rather than absorption by cool outÑow.
5.4. Ha Emission : Magnetospheric Origin?In our picture, the Ha red emission is due to accretion
and is produced in the magnetospheric Ñow, while thebroad weak symmetric outer wings of Ha may be due tomagnetic turbulence. The outer wings are very well Ðtted bya Gaussian centered at the stellar rest frame, which stronglysuggests a turbulent origin, and they correlate well with theHa redshifted emission, indicating a magnetospheric origin.Muzerolle, Calvet, & Hartman (1998a) suggest that Starkbroadening e†ects are responsible for the extended Hawings. However, in the case of DR Tau, it is hard to imaginethat the far wings of Ha are the result of Stark broadeningof a strong centered emission component since most of theHa emission is obviously redshifted and there is no indica-tion of such a component in most of the higher Balmer lines
where the inÑuence of wind absorption should be weaker.The higher Balmer lines do not generally show this broadcomponent, indicating it is optically thin there. If Starkbroadening is important in the formation of the Balmerlines, its e†ects should be testable using the high-lyingmembers of the series where Stark broadening is morenoticeable. Such comparisons have clearly shown the role ofStark broadening in solar Ñares, permitting a determinationof the local electron density in the line formation region(e.g., Johns-Krull et al. 1997, and references therein). Inaddition, Ardila et al. (2001) Ðnd that the Mg II h ] k lineshave proÐles very similar to Ha ; this would be very sur-prising if Stark broadening were a major contributor.
The variable blue emission peak may get a contributionfrom the far side of the magnetospheric Ñow but may alsopartially arise in a sporadic, jetlike outÑow launched rela-tively close to the star. Goodson, & Winglee (1999)Bo� hm,presented a model for the formation of jets and outÑowsfrom young stellar objects that could explain the variableblue emissions we observe in DR Tau. They show that dif-ferential rotation between the star and the disk twists theclosed magnetic Ðeld lines, leading to their expansion.While inÑating they drive plasma in an outÑow that is setfree when the expanded Ðeld lines reconnect. As long as themagnetically coupled regions of the star and disk rotatedi†erentially with respect to each other, this process canstart again and new blueshifted emissions will be produced.
Observationally, the velocity of the Ha blue emissionpeak is comparable with the terminal outÑow velocity mea-sured by the position of the blue emission peak of the [O I]j6300 line. This forbidden line supposedly forms manystellar radii from the star and in DR Tau is commonlyfound in the range [160 km s~1\ v\ [120 km s~1(Appenzeller et al. 1988 ; Hartigan et al. 1995). The red andblue emissions of Ha do not seem to be related to each otherin terms of their variations, as can be seen in Figure 9, wherethere is no sign of correlation between them, so they maynot originate in the same regions. Most of the variability inHa is on the blue side, showing that the source of blueemission is much more ephemeral. The intensities of theblue and red emissions are also very di†erent from Ha toHd (see Fig. 2 in Appenzeller et al. 1988). While the reddecreases drastically, the blue remains almost constant. WeveriÐed this behavior in our data. All this points to anoutÑow origin for the blueshifted emission.
5.5. Magnetospheric Toy Model for HaAlthough we picture the redshifted emission as arising
from the accretion Ñow, the observed Ha proÐles do notlook at all like the published theoretical ones at any inclina-tion. As a result, we decided to investigate the inÑuence ofthe many geometrical parameters and the Ñux distributionalong the Ðeld lines on the resulting proÐle characteristics.We computed a very simple dipole model where materialfree-falls from the disk to the star along the Ðeld lines. Thedisk was considered to be opaque and emissionless at allfrequencies considered here (Ha), and we did not take intoaccount the emission from a hot accretion ring or spot onthe stellar surface. We calculated the velocity vector andattributed a Ñux value to every point on the dipole stream-lines. We had the option of adding a turbulent velocitycomponent with a random Gaussian distribution to eachvelocity point. We then calculated the projected velocities inthe observerÏs direction and added up the Ñux of the points
No. 6, 2001 SPECTRAL VARIABILITY OF DR TAURI 3357
FIG. 11.ÈTheoretical line proÐle computed with a Ñux similar to theone used in Hartmann, Hewett, & Calvet (1994), inclination\ 10¡, R
*\
and period \ 4 days.2.6 R_
, M*
\ 0.4 M_
,
that fell in the same 10 km s~1 velocity bin. The intensity ofthis imaginary line proÐle was chosen to match the Haobserved intensities and the shape of the proÐle reÑects thegeometry of the system and the line emissivity in the magne-tosphere.
Using this simpliÐed toy model, we nevertheless obtainline proÐles that demonstrate the same general character-istics as those of Hartmann, Hewett, & Calvet (1994) exceptfor the occasional redshifted absorption, which appearswhen material in the funnel is projected along our line ofsight in front of a hot ring (or spot) that is not included inour model. In Figure 11 we show an example obtained witha source function similar to the one used by Hartmann et al.(1994) along the Ðeld lines and a similarly low inclination.The shape of our proÐle is nearly identical to that of Hart-mann et al. (1994 ; their Fig. 7). This shape is very di†erentfrom the observed Balmer lines of DR Tau, so we decided touse di†erent line emissivity functions along a streamline.The resulting proÐle with a line emissivity that stronglyincreases near the star is presented in Figure 12. This theo-retical proÐle supports the decomposition we adopted sinceit shows that a strongly redshifted emission component canbe produced in a magnetospheric accretion Ñow evenwithout a wind absorbing in the blue wing. We obtain pro-Ðles strongly peaked in the red for 0¡ \ i\ 20¡, but forhigher inclinations the proÐles no longer match theobserved Balmer lines. Assuming a longer rotation period(which requires a larger magnetosphere, assuming all othersystem parameters stay Ðxed to maintain equilibrium) shiftsthe red emission peak toward higher velocities (P\ 6 dayspeaks at D150 km s~1, and P\ 9 days peaks at D165 kms~1). We would ascribe the appearance of the occasionalstrong blue peak to an additional high-velocity outburst, asdiscussed earlier. These outbursts are responsible for theblueshifted peak in the variance proÐle in Figure 7, alongwith the cool wind absorption component (the two aremore distinct in the Hb variance proÐle). The Ñux in thefunnel is the speciÐed Ñux function (which typically dependson the position of the material along the accretion funnel)divided by the stellar continuum Ñux at Ha. We can thenestimate the temperature required to give this Ñux function
FIG. 12.ÈTheoretical line proÐle computed with power-law Ñux(Pr~7), inclination\ 10¡, and period \ 4R
*\ 2 R
_, M
*\ 0.4 M
_,
days. The proÐle was smoothed with a window width of 3.
by assuming optically thick emission in LTE. Thus, if theÑux function is set to 1 everywhere, the temperature alongthe funnel is constant and equal to the stellar temperature.While LTE is certainly not a proper assumption for thefunnel Ñow, it permits a general estimate of the tem-peratures required to produce the modeled emission. Com-paring our results in Figure 11 with the parameters ofHartmann et al. (1994) suggests that such an estimate isgood to a factor of 2 or better. An estimate of the tem-perature structure that produced the proÐle in Figure 12 isshown in Figure 13, together with the Ñux function used toobtain the line proÐle of Figure 12.
Figures 12 and 13 show that it is possible to obtain lineproÐles that look like the Balmer lines observed in DR Tauwith a heuristically constructed magnetospheric accretionconÐguration and a reasonable temperature structure. Wedo note that the CTTSs T Tau and TW Hya are alsothought to be viewed nearly pole-on, as we advocate herefor DR Tau. The Ha line proÐles of these two stars peak
FIG. 13.ÈFlux function and temperature corresponding to the modelin Fig. 14.
3358 ALENCAR, JOHNS-KRULL, & BASRI Vol. 122
much closer to the stellar rest velocity and are generallymore symmetric than that of DR Tau. These di†erencesmust be accounted for in any complete theory of CTTSphenomena. One strong di†erence between these stars istheir mass accretion rates. Calvet & Gullbring (1998) derivean accretion luminosity of 2.74 for DR Tau, implying aL
_mass accretion rate of a few times 10~7 yr~1, whileM_Muzerolle et al. (2000) determine a mass accretion rate of
5 ] 10~10 yr~1 for TW Hya. If the mass accretion itselfM_has an inÑuence on the temperature structure in the funnel
Ñow, such a tremendous di†erence in accretion rates couldaccount for the observed di†erence between stars such asDR Tau compared with TW Hya and T Tau. We note thatthere is no complete theory of the heating of the funnel Ñow.The investigations by Martin (1996) determine funnel Ñowtemperature that are too low to account for the observedBalmer emission. As a result, studies such as those of Hart-mann et al. (1994), Muzerolle et al. (1998a), and the modelspresented above rely on purely ad hoc temperature proÐles.
5.6. He I L inesBecause of the high temperatures and densities required
to form the He I line in emission, it is probably produced inthe accretion Ñow close to the stellar surface or the accre-tion shock itself. Since it correlates very well with the Hared emission peak, we have a self-consistent picture, whichstrengthens our assertion that the red emission peak of theHa lines forms close to the stellar surface. The same kind ofcorrelation is found between the Hb red emission peak andthe He I line. This conÐrms the trend shown by Alencar &Basri (2000) between the Ha, Hb, and He I veiling-correctedequivalent widths of several CTTSs.
The narrow and broad He I components are expected tocorrelate with the veiling, the former because it is enhancedby accretion and the latter because it is thought to arise inthe accretion Ñow or shock. However the equivalent widthsof the He I components, either as observed or corrected forveiling, generally do not show a correlation with the veilingitself. A slight correlation is seen between the veiling-corrected BC and the veiling itself, but nothing is apparentbetween the NC or the redshifted absorption componentand the veiling. While this result is difficult to understand interms of the magnetospheric accretion model, we must keepin mind that our veiling values are not extremely reliable.Additional studies with higher quality data are needed toconÐrm or refute this Ðnding.
As pointed out by Hartmann et al. (1994) the He I linelooks very similar to the theoretical magnetospheric accre-tion proÐles. Since the He I line does not have any windcontribution and is not expected to have any signiÐcantStark broadening, it is a good candidate to compare withthe theoretical proÐles. We do not have these proÐles spe-ciÐcally calculated for the He I line but decided to use scaledHartmann et al. (1994) Hb proÐles (kindly provided by L.Hartmann) in the comparison instead. While the He I
source function will not look exactly like the Hb one, thetemperature structure in the magnetospheric Ñow is stillpoorly understood and is currently speciÐed in an ad hocfashion. Our primary assumption is then that the shape ofthe He I theoretical proÐles should be similar to those thathave been calculated for Hb.
In Figure 14a we show a typical He I line proÐle with ascaled theoretical proÐle overplotted. We have added to thetheoretical proÐle our Gaussian Ðt of the observed spec-
FIG. 14.ÈHe I line proÐle decomposition and comparison with thetheory. The solid line is the observed proÐle. In (a) the dashed line is theproÐle Ðt obtained by adding the two components shown in (b). The dash-dotted line is the theoretical proÐle with the Ðtted NC added to it. In (c) weshow only the BC. The solid line is the observed proÐle with the NCsubtracted, the dashed line is the BC Ðt, and the dash-dotted line is thetheoretical proÐle.
trumÏs NC. The theoretical proÐles should be comparedonly with the BC and the occasional redshifted absorptionsince the NC, despite being a†ected by the accretionprocess, is not thought to be produced in the magneto-spheric Ñow itself. We note that, although the entire He I
line may look asymmetric, the individual components arevery symmetric (Fig. 14b). The fact that the BC is normallyfound blueshifted while the NC is not gives a false impres-sion of asymmetry to the overall line shape. The theoreticalproÐle in Figure 14 corresponds to an inclined magneto-spheric model with i \ 60¡ (chosen as the best match to theshape of the observed He I line) and hot ring emission at thebase of the magnetosphere (Fig. 9 in Hartmann et al. 1994).In general, the published theoretical proÐles tend to be veryasymmetric, with a lack of redshifted low-velocity materialdue to occultation by the disk of the material moving slowlyaway from the observer. Extended red wings are also com-monly produced at lower inclinations or with di†erent geo-metric assumptions (e.g., hot spot emission instead of aring). In this case the lack of blueshifted emission in the farwings is caused by occultation by the star of the high-velocity material moving toward the observer. We do notnormally see any of these asymmetries in our sample.
An exceptionally asymmetric He I proÐle is shown inFigure 15. This time the agreement with the theoreticalproÐle shape is much better, but unfortunately proÐles likethis one are quite rare in our sample for DR Tau, and againthe inclination of the model proÐle is much di†erent thanthat assumed by Kenyon et al. (1994) or argued for here.The theoretical proÐles tend to be either blueshifted andasymmetric (higher inclinations) or centered and symmetric(lower inclinations). They are not blueshifted and symmetriclike most of the He I BCs observed here. In Hartmann et al.
No. 6, 2001 SPECTRAL VARIABILITY OF DR TAURI 3359
FIG. 15.ÈAsymmetric He I line proÐle decomposition and comparisonwith the theory. The lines have the same meaning as in Fig. 11 except thatnow there is a redshifted absorption present.
(1994) the case that comes closest to what we observed is alarge magnetosphere, which favors the material starting outin the disk near corotation radius. The proÐles look quitesymmetric even at higher inclinations, but it is unclear thatsuch proÐles would also appear blueshifted. Above, we haveargued for temperature structure in the magnetosphere thatstrongly weights the Ha emissivity toward the stellarsurface. It is likely that such a structure would result instrongly redshifted He I lines from the magnetosphere aswell. Taken together, these results seem to point to a non-funnel origin for most of the He I broad component.
Beristain et al. (2001) noted that the He I BC of high massaccretion rate CTTSs, like DR Tau, tend to be blueshiftedand attributed this to a possible line formation in a hotwind. We suggest that this ““ hot wind ÏÏ is actually theunstable outbursts found by Goodson et al. (1999) andHayashi, Shibata, & Matsumoto (1996). These result fromthe fact that changes in both the disk material and stellar
magnetic Ðeld conÐguration make it difficult for the star-disk interface to stay in equilibrium.
Another possibility is that the He I and Ca II broad emis-sion components are mainly produced by magnetosphericturbulence. Typical values of magnetic Ðelds measured inCTTSs (D2.5 kG; Johns-Krull & Valenti 2000) and theo-retical gas densities used in magnetospheric accretionmodels (D1 ] 10~9 g cm~3) yield velocities of D200Alfve� nkm s~1. It is not clear that turbulence can be generated upto this velocity ; however, with accreting gas impacting thestellar surface at velocities equal to or greater than this,such large values of magnetic turbulence may result. Thisplausible turbulent origin would help explain the symmetryof these components.
6. CONCLUSIONS
We analyzed more than 100 spectra of the classical TTauri star DR Tau. We showed that the emission-line pro-Ðles vary constantly, but not with a unique period, implyingthat the variations are caused by nonsteady accretion andoutÑow processes. Many di†erent emission-line changes arecorrelated, showing that di†erent line-forming regions area†ected by the same events. The published theoretical mag-netospheric accretion proÐles are very di†erent from ourobserved ones. We propose for DR Tau that the very sym-metric Ca II and He I broad components may be producedby magnetospheric turbulence while the Balmer lines orig-inate in the magnetosphere with a line emissivity thatincreases strongly near the star. We computed simple dipolemodels and showed that the observed redshifted BalmerproÐles can be produced in such a magnetospheric accre-tion conÐguration observed at low inclination. We also notethe presence of sporadic outÑow events that can be seen inthe Balmer and He I lines.
This research is based on data collected on the Shane 3 mand CAT telescopes at Lick Observatory run by the Uni-versity of California. We would like to thank AnthonyMisch, Claude Bertout, Natalie Stout-Batalha, and CelsoBatalha, who helped gather the proÐles presented here.S. H. P. A. acknowledges support from the ConselhoNacional de Desenvolvimento e andCient•� Ðco Tecnolo� gicoFAPESP grant 2000/06244-9-Brazil. C. M. J.-K. is pleasedto acknowledge partial support from NASA grant NAG5-8209.
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