The T Tauri Sun

6.3 The T Tauri Sun

6.3.1 The magnetic field of the T Tauri Sun

Surface magnetic fields have been successfully measured on “solar-analog” T Tauri stars, using Zeeman broadening (Johns-Krull et al., 1999 ; Valenti and Johns-Krull, 2004 ; Johns-Krull, 2007a ). For example, BP Tau (a 0.65 $M ⊙$ CTTS) maintains a mean magnetic field strength ( $∑$ Bf) of 2.6 $±$ 0.3 kG, i.e., the equivalent of sunspots for f = 1. Such field strengths exceed the equipartition value of T Tauri photospheres ( $≈$ 1 kG, Johns-Krull et al. 1999 ; Johns-Krull 2007a ), which may be a consequence of near-total surface filling (Solanki, 1994 ; Johns-Krull, 2007a ). T Tauri photospheres are therefore dominated by magnetic pressure rather than thermal pressure, in contrast to MS stars and the Sun but similar to stellar coronae in general. The average surface magnetic fields of CTTS also exceed a prediction from X-rays, based on a correlation between X-ray luminosity and photospheric magnetic flux valid for MS stars and the Sun (Pevtsov et al., 2003 ). Upper limits of the net polarization of photospheric lines suggest that the photospheric magnetic fields form predominantly in small-scale structures, although a dipole may dominate at large distances (Johns-Krull et al., 1999 ; Valenti and Johns-Krull, 2004 ; Johns-Krull, 2007a ). Large dipole components are also suggested from observations of spots concentrated at the poles of T Tauri stars, similar to active MS solar analogs (e.g., Joncour et al., 1994 ; Hatzes, 1995 ; Rice and Strassmeier, 1996 ). In any case, the observed fields should be strong enough to truncate circumstellar disks indeed (Johns-Krull et al., 1999 ; Johns-Krull, 2007a ), even if rather complex (non-dipolar) surface magnetic field distributions are assumed (Gregory et al., 2006 ).

Zeeman–Doppler Imaging techniques have successfully been applied to extremely active solar analogs in the PMS phase. A particularly clear case was presented by Donati et al. (2000 ), finding solar-like differential rotation on a post-T Tauri star (see also Section 4.3.1). More recently, Donati et al. (2007 ) have reconstructed a rather complex, large-scale magnetic topology on the CTTS V2129 Oph; they found a relatively weak dipole but stronger octupolar fields that are tilted against the rotation axis, with strong near-polar spots. The accretion footpoints are also found to be located at high latitudes. An attempt was made at extrapolating the fields to the inner rim of the disk, showing that some field lines should successfully accrete toward the observed hot spots.

At coronal levels, X-ray rotational modulation provides information on the large-scale distribution of stellar magnetic fields. Rotational modulation is widespread among extremely active T Tauri stars (Flaccomio et al., 2005 ). Some 10% of the studied stars in the Orion X-ray sample show such evidence, suggesting that: i) the X-ray emitting active regions are not homogeneously distributed on the surface, i.e., despite the X-ray saturation level reached by these stars, the surface cannot be filled with X-ray-bright magnetic loops; ii) the X-ray emitting regions responsible for the rotational modulation are directly associated with the surface and cannot extend much beyond $R ∗$ (Flaccomio et al., 2005 ). A comparison of the modulation depth with the Sun’s modulation in fact shows that the longitudinal inhomogeneities are similar (Flaccomio et al., 2005 ).

6.3.2 The ultraviolet T Tauri Sun

A defining property of (accreting) classical T Tauri stars is their strong line emission of, e.g., H $α$ or Ca ii H & K. These strong lines were initially thought to be evidence of massive chromospheres similar to those seen on the Sun or in cool stars (see review by Bertout 1989 ), and the discovery of strong UV lines such as those of Si ii, Si iv, and C iv – equivalent to “transition region” lines in the Sun formed above 10⁴ K – supported this picture.

However, when compared with MS stars, including chromospherically very active examples, UV line and continuum emission is up to 10²–10⁴ times stronger in CTTS (see example of TW Hya in Table 4; Canuto et al. 1982 , 1983 ; Bouvier 1990 ; Valenti et al. 2000 ), regardless of the photospheric effective temperature or the stellar rotation period but correlated with the mass accretion rates derived from optical continuum data (Bouvier 1990 ; Johns-Krull et al. 2000 ; Figure 32 a below). Further, UV or H $α$ line surface fluxes of CTTS show, in contrast to more evolved stars, no correlation with coronal X-rays, the latter being in the range of RS CVn-type active binary systems or very active MS stars but the UV/H $α$ lines showing a wide range of excess flux (Bouvier, 1990 ).

Coronal and “chromospheric/transition region” fluxes are thus not correlated in CTTS, contrasting strongly with MS and subgiant stars for which a sharp correlation is taken as evidence for a common physical heating mechanism (operating in related magnetic fields; Section 5.6, Figure 22 ). An additional mechanism must be responsible for the optical/UV line flux excess. Apart from the line excess fluxes, there is also a strong blue continuum excess that leads to “veiling” in the optical spectrum, i.e., a filling-in of absorption lines by continuum emission; this emission is also not compatible with chromospheric radiation. The most obvious property common only to CTTS among the stars considered above is accretion; downfalling material could provide the energy to generate the optical/UV excess (Bertout et al., 1988 ; Basri and Bertout, 1989 ).

Nearly free-falling gas can be heated to maximum temperatures

in shocks forming at the bottom of magnetic accretion funnels (Calvet and Gullbring, 1998 ). UV and optical line emission could thus provide diagnostics for the accretion velocity, the mass accretion rate, and possibly the surface filling factor of accretion funnels.

The present consensus, based on such concepts as well as line profile properties and correlations with the mass accretion rate, is that the UV excess emission originates from material heated in accretion shocks (e.g., Calvet and Gullbring, 1998 ; Gullbring et al., 1998 ). Some of the emission lines (e.g., H $α$ , Ca ii) may also form in the accretion funnels themselves, or in stellar winds (Ardila et al., 2002 ).

6.3.3 The X-ray T Tauri Sun in time

Feigelson et al. (2002b ), Wolk et al. (2005 ), and Telleschi et al. (2007b ) presented X-ray studies of near-solar-mass stars (stars in the ranges of $0.7 M ⊙ ≤ M ≤ 1.4M ⊙$ and $0.9M ⊙ ≤ M ≤ 1.2 M ⊙$ , respectively, in the former two studies of the Orion Nebula cluster, and wider in the latter study of the Taurus star-forming region). The sample ages typically comprise the log t = 5.5–7 range and contain both disk-surrounded and disk-less T Tauri stars. The median X-ray luminosity in the Orion sample is found at log L_X = 30.25, i.e., three orders of magnitude above the average solar X-ray output, but there is evidence for a slow decay with age, log L_X $∝$ t^–1.1 (Feigelson et al., 2002b ; Wolk et al., 2005 ). A shallower decay was reported by Preibisch and Feigelson (2005 ) for the same stellar cluster, with an exponent between –0.2 and –0.5, but when considering normalized L_X / L_bol or average surface X-ray flux, then both Feigelson et al.’s and Preibisch & Feigelson’s studies indicate that the L_X decay law is roughly compatible with full saturation (i.e., L_X / L_bol $≈$ 10^–3) as the star descends the Hayashi track and its bolometric luminosity is decreasing (Feigelson et al., 2002b ). Telleschi et al. (2007b ) used the Taurus sample over a wider mass range but removed the strong L_X versus mass correlation in order to normalize the X-ray evolutionary behavior to a solar-mass star. The slope of the L_X vs. age correlation is fully compatible with the Orion results, with a power-law index of –0.36 $±$ 0.11 although the correlation is dominated by scatter from other sources, and its significance is marginal.

The evolutionary L_X decay is thus qualitatively different from that in MS stars: it is due to stellar contraction (and perhaps a change in the internal dynamo while the star transforms from a fully convective to a convective-radiative interior); in contrast, the decay of L_X in MS stars is due to stellar spin-down while the stellar structure and size remain nearly constant. Figure 31 shows the long-term evolution of the median X-ray output from PMS stages to the end of the MS evolution, for G-type stars with ages $>$ 10 Myr and K-type stars with ages $<$ 10 Myr (because the predecessors of MS G stars are PMS K stars; data from Güdel 2004 ). The slight trend toward decreasing L_X at ages $<$ 10 Myr follows approximately L_X $∝$ t^–0.3, in agreement with the individual trends for the Orion and the Taurus samples, albeit the scatter is large. No “onset” of activity can be seen back to ages $<$ 1 Myr.

Figure 31:

Evolution of the X-ray luminosity of a near-solar mass star. A sample of field stars from the “Sun in Time” program is shown by filled circles. Open circles show the median L_X for G-type stars in open clusters for ages $>$ 10 Myr, and for K-type stars in open clusters and star-forming regions for ages $<$ 10 Myr. The error bars show the approximate 1 $σ$ scatter in the samples (adapted from Güdel, 2004

, original data from references given therein; © Springer Verlag, reprinted with permission).

In summary, the age evolution of the X-ray output is modest in PMS stars, the bulk of the X-ray output being determined by other stellar properties. There are at least four such parameters that have been discussed in the recent literature: bolometric luminosity, mass, rotation, and mass accretion rate. I briefly summarize these parameter dependencies in turn:

Similar to young MS stars, CTTS and WTTS straddle the empirical X-ray saturation limit, i.e., L_X $≈$ 10^–3.5 L_bol (e.g., Preibisch et al., 2005 ; Telleschi et al., 2007b ), pointing to a dynamo process somehow related to the dynamo in MS stars. The long-term evolution of L_bol may then be the principal parameter for the long-term evolution of L_X.
In contrast to MS stars, PMS stars show a correlation between L_X and mass (e.g., Feigelson et al., 1993 ; Preibisch et al., 2005 ; Telleschi et al., 2007b ). However, if one restricts the MS sample to saturated, young stars, then they follow such a relation as well, owing to the well-known mass-L_bol relation on the MS, L_bol $∝$ M³. Conversely, the flatter L_bol-M relation for a given PMS isochrone combined with the saturation law yields an L_X vs. mass correlation that is compatible with the observed relation (Telleschi et al., 2007b ), i.e., the two relations are interdependent.
Rotation is one of the main drivers of the magnetic dynamo and therefore of magnetic activity in MS stars. The near-saturation state of most PMS stars, however, suggests that rotation cannot be a key parameter. This is borne out by explicit correlation studies of the Orion (Preibisch et al., 2005 ) and the Taurus (Briggs et al., 2007 ) samples of T Tauri stars that show little in the way of a correlation as seen in MS stars. On the contrary, some apparent trends in this direction are the result of population bias (Briggs et al., 2007 ). The absence of a decrease in the L_X / L_bol ratio up to rotation periods of at least 10 d (in contrast to MS stars) can be explained by the convective turnover time in PMS stars being much larger, yielding smaller Rossby numbers for a given rotation period and therefore saturation up to longer rotation periods (Preibisch et al., 2005 ; Briggs et al., 2007 ).
Accretion could influence L_X either by generating X-rays itself, or by suppressing coronal X-ray production. Empirically, L_X is correlated with the mass accretion rate $˙ M$ but this is not a physical correlation for the following reason. Muzerolle et al. (2003 ), Muzerolle et al. (2005 ), and Calvet et al. (2004 ) described a relation between $M˙$ and the stellar mass $M$ , approximately reading $M˙ ∝ M 2$ . Because L_X correlates with L_bol and therefore, for a typical age isochrone, with mass, L_X must correlate with $M˙$ . Telleschi et al. (2007b ) removed this trend from the Taurus sample. Although the two parameters still do not seem to be fully independent, a clear correlation cannot be established. Accretion does seem to influence magnetic energy output in more subtle ways, however, to be discussed in the Section 6.3.4.

6.3.4 Coronal excesses and deficits induced by activity?

If the photoelectric absorption by the accreting gas is small, then the softest X-ray range may reveal the high-temperature tail of the shock emission measure thought to be responsible for the UV excesses (Section 6.3.2). Telleschi et al. (2007b ) and Güdel et al. (2007b ) identified an excess in the O vii/O viii Ly $α$ flux (or luminosity) ratio in CTTS when compared with WTTS or MS stars, the so-called X-ray soft excess of CTTS (Figure 32 b). In the most extreme case of the CTTS T Tau, the excess O vii flux is such that this line triplet, formed at only $≈$ 2 MK, is the strongest in the soft X-ray spectrum (see Figure 33 ; Güdel and Telleschi 2007 ).

Figure 32:

Left (a): The ultraviolet excess of CTTS. The figure shows the mean continuum surface flux at $≈$ 1958 Å vs. the stellar effective temperature. CTTS are shown by the solid circles, naked (diskless) TTS by asterisks, and MS stars by squares. The solid lines define a fit to the MS stars, the dashed lines giving the lower and upper bounds (from Johns-Krull et al., 2000

, reproduced by permission of AAS). Right (b): The X-ray soft excess in CTTS. The figure shows the ratio between O vii r and O viii Ly $α$ luminosities (each in units of erg s^–1) vs. L_X. Labels give initial letters of stellar names. Crosses mark MS stars (from Ness et al., 2004 ), triangles solar analogs (from Telleschi et al., 2005

), filled (red) circles CTTS, and open (blue) circles WTTS. Two solutions (high- and low-absorption) are given for T Tau, connected by a dotted line. The solid line is a power-law fit to the MS stars with L_X $>$ 10²⁷ erg s^–1 (from Güdel and Telleschi, 2007

, reprinted with permission).

Figure 33:

Comparison of fluxed X-ray photon spectra of (from top to bottom, all from XMM-Newton RGS) the active binary HR 1099, the WTTS V410 Tau, the CTTS T Tau, the T Tau spectrum modeled after removal of absorption (two versions, using N_H = 3 $×$ 10²¹ cm^–2 and 4.9 $×$ 10²¹ cm^–2, respectively), and the inactive MS star Procyon. The bins are equidistant in wavelength (from top to bottom, the bin widths are, respectively, 0.025 Å, 0.070 Å, 0.058 Å, 0.058 Å, 0.058 Å, and 0.010 Å; from Güdel and Telleschi 2007

, reprinted with permission).

Interestingly, however, the X-ray soft excess in CTTS is comparatively moderate, the L(O vii r) / L(O viii) line flux ratio being enhanced by factors of typically $≈$ 3–4 times that of equivalent MS stars or WTTS (Figure 32 ). Furthermore, the excess X-ray line fluxes do not seem to be correlated with the UV line excesses but are correlated with the overall stellar coronal activity level traced, for example, by the O viii Ly $α$ line flux (Güdel and Telleschi, 2007 ). It appears that the X-ray soft excess depends on the level of magnetic (“coronal”) activity although it is, at the same time, related to the presence of accretion. The two dependencies may point to an interaction between accretion and magnetic activity at “coronal” heights.

The shock interpretation of the softest X-rays and the X-ray soft excess is appealing, but remains controversial until a larger sample of CTTS with various accretion properties has been interpreted. In particular, given the appreciable accretion rates, high shock densities of order 10¹²–10¹⁴ cm^–3 are expected, as first indeed reported from density-sensitive line diagnostics of O vii and Ne ix in the CTTS TW Hya, forming at only a few MK (Kastner et al., 2002 ; Stelzer and Schmitt, 2004 ). However, some accreting PMS stars show much lower densities, such as AB Aur (Telleschi et al., 2007b ) and T Tau (Güdel et al., 2007b ); the same discrepancy between expected and observed densities has also been reported from ultraviolet density diagnostics (Johns-Krull et al., 2000 ).

In stark contrast to the X-ray soft excess and the UV excess described above, it is now well established that CTTS show a moderate suppression of 0.1–10 keV soft X-ray emission, typically by a factor of $≈$ 2 when compared with WTTS of similar properties (Strom and Strom, 1994 ; Damiani et al., 1995 ; Neuhäuser et al., 1995 ; Stelzer and Neuhäuser, 2001 ; Flaccomio et al., 2003 ; Preibisch et al., 2005 ; Telleschi et al., 2007a ). Although selection/detection bias or different photoelectric absorption has been quoted to be responsible for these differences (see Güdel 2004 and references therein), the luminosity deficit in CTTS is now thought to be real; it appears that accretion suppresses coronal heating in a fraction of the coronal volume (Preibisch et al., 2005 ), or at least leads to larger amounts of cooler plasma, which is perhaps the same plasma inferred from the X-ray soft excess (Telleschi et al., 2007b ; Güdel and Telleschi, 2007 ). Alternatively, the presence of a circumstellar disk could strip the outer parts of the stellar corona, thus reducing L_X (Jardine et al., 2006 ).

6.3.5 X-ray flaring of the T Tauri Sun

A high level of near-continuous flaring is found in PMS solar-mass stars. As much as half of the emitted X-ray energy, if not more, may be due to strong flares (Montmerle et al., 1983 ), and many TTS are nearly continuously variable probably also owing to flares (Mamajek et al., 2000 ; Feigelson et al., 2002a ; Preibisch and Zinnecker, 2002 ; Skinner et al., 2003 ). Examples with extreme luminosities and temperatures up to 100 MK have been reported (e.g., Preibisch et al., 1995 ; Skinner et al., 1997 ; Tsuboi et al., 1998 , 2000 ; Imanishi et al., 2002 ). The most extreme flares are found on CTTS and protostars, a possible hint at star-disk magnetic interactions during flares. Wolk et al. (2005 ) studied frequency and properties of flares in the Orion Nebula cluster, concluding that the median peak luminosity of their sample was log L_X = 30.97, with extremely hard spectra at peak time. The median electron temperature was found at 7 keV. An analogous study has been presented by Stelzer et al. (2007 ) for T Tauri stars in the Taurus Molecular Cloud. The extreme flaring recorded on these PMS stars may have an important bearing on coronal heating (see Section 5.8) and on the alteration of solids in the young stellar environment (see Section 6.6).

6.3.6 The radio T Tauri Sun in time

Early VLA surveys quickly reported strong radio emission from both CTTS and WTTS. Somewhat unexpectedly, however, radio emission comes in two principal flavors. The early, pioneering studies by Cohen et al. (1982 ), Bieging et al. (1984 ), Cohen and Bieging (1986 ), Schwartz et al. (1984 ), and Schwartz et al. (1986 ) recognized thermal wind-type emission with rising spectra and in cases large angular sizes for several CTTS. This radio emission can then be used to estimate mass loss rates; these are found to range up to $∼<$ 10^–7 $M ⊙$ yr^–1 (André et al., 1987 ). The partly enormous kinetic wind energy derived under the assumption of a uniform spherical wind suggests anisotropic outflows while structural changes in the radio sources indicate variable outflows, probably along jet-like features; at shorter radio wavelength, dust emission from the disk becomes apparent as well (Cohen and Bieging, 1986 ; Rodríguez et al., 1992 , 1994 ; Wilner et al., 1996 ). The thermal radio emission tells us nothing about the presence or absence of stellar magnetic fields. As described earlier, CTTS do show many signatures of magnetic activity, but whatever the possible accompanying radio emission, it seems to be absorbed by the circumstellar ionized wind.

The situation is different in WTTS in which the presence of huge flares (Feigelson and Montmerle, 1985 ; Stine et al., 1988 ; Stine and O’Neal, 1998 ), longer-term variability, and falling spectra clearly point to non-thermal gyrosynchrotron emission (Bieging et al., 1984 ; Kutner et al., 1986 ; Bieging and Cohen, 1989 ; White et al., 1992a ; Felli et al., 1993 ; Phillips et al., 1996 ) analogous to radio emission observed in more evolved active stars. Conclusive radio evidence for the presence of solar-like magnetic fields in WTTS came with the detection of weak circular polarization during flares but also in quiescence (White et al., 1992b ; André et al., 1992 ; Skinner, 1993 ). Extremely energetic particles radiating synchrotron emission may be involved, giving rise to linear polarization in flares on the WTT star HD 283447 (Phillips et al., 1996 ). VLBI observations showing large ( $∼ 10 R ∗$ ) magnetospheric structures with brightness temperatures up to T_b $≈$ 10⁹ K fully support the non-thermal picture (Phillips et al., 1991 ).

As a WTT star ages, its radio emission drops rapidly on time scales of a few million years from luminosities as high as 10¹⁸ erg s^–1 Hz^–1 to values around or below 10¹⁵ erg s^–1 Hz^–1 at ages beyond 10 Myr. Young age of a star is thus favorable to strong radio emission (O’Neal et al., 1990 ; White et al., 1992a ; Chiang et al., 1996 ), whereas toward the subsequent ZAMS stage it is only the very rapid rotators that keep producing radio emission at the 10¹⁵ erg s^–1 Hz^–1 level (Carkner et al., 1997 ; Magazzù et al., 1999 ; Mamajek et al., 1999 ).

6.3.7 The composition of the T Tauri Sun’s corona

Initial studies of a few accreting T Tau stars, in particular the old ( $≈$ 10 Myr) TW Hya, have shown an abundance pattern in the X-ray source similar to the IFIP effect although the Ne/Fe abundance ratio is unusually high, of order 10 with respect to the solar photospheric ratio, and the N/O and N/Fe ratios are enhanced by a factor of $≈$ 3.

These anomalous abundance ratios have been suggested (Stelzer and Schmitt, 2004 ; Drake and Testa, 2005 ) to reflect depletion of Fe and O in the accretion disk where almost all elements condense into grains except for N (Savage and Sembach, 1996 ; Charnley, 1997 ) and Ne (Frisch and Slavin, 2003 ) that remain in the gas phase which is accreted onto the star. If accretion occurs predominantly from the gas phase in the higher layers of the disk while the grains grow and/or settle at the disk midplane, then the observed abundance anomaly may be a consequence.

Larger systematics have made this picture less clear, however. Several CTTS and WTTS have revealed large Ne/Fe ratios ( $≈$ 4 or higher), much larger than in MS active solar analogs (Kastner et al., 2004 ; Argiroffi et al., 2005 , 2007 ; Telleschi et al., 2005 , 2007b ; Günther et al., 2006 ) but similar to RS CVn binaries (Audard et al., 2003b ). In contrast, the CTTS SU Aur reveals a low Ne/Fe abundance ratio of order unity (Robrade and Schmitt, 2006 ; Telleschi et al., 2007b ), similar to some other massive CTTS (Telleschi et al., 2007b ).

Partial clarification of the systematics has been presented by Telleschi et al. (2007b ) (see also Güdel et al., 2007b ) who found that

the abundance trends, and in particular the Ne/Fe abundance ratios, do not depend on the accretion status but seem to depend on spectral type or surface T_eff, the later-type stars showing a stronger IFIP effect (larger Ne/Fe abundance ratios); see Figure 34 .
the same trend is also seen in disk-less ZAMS stars.

Anomalously high Ne/O abundance ratios remain, however, for TW Hya (Stelzer and Schmitt, 2004 ) and V4046 Sgr (Günther et al., 2006 ) when compared to the typical level seen in magnetically active stars, including PMS objects. The initial idea proposed by Drake et al. (2005 ) was that the selective removal of some elements from the accretion streams should occur only in old accretion disks such as that of TW Hya where cogulation of dust to larger bodies is ongoing, whereas younger T Tauri stars still accrete the entire gas and dust phase of the inner disk. However, the old CTTS MP Mus does not show any anomaly in the Ne/O abundance ratio (Argiroffi et al., 2007 ). Larger samples are needed for clarification.

Figure 34:

The coronal Fe/Ne abundance ratio (with respect to the solar photospheric mixture) in the X-ray sources of various PMS stars and very active MS stars as a function of spectral class. Symbols mark different types of stars: squares: active MS stars; triangles: WTTS; diamonds: CTTS. For details and references, see Telleschi et al. (2007b ) and Güdel et al. (2007b ).