eventually?
)
The history of magnetic field measurements in cool and, in particular, in sun-like stars, is not easily followed. The fundamental paradigm of magnetic fields leading to chromospheric and coronal emission, as observed on the Sun, has motivated clear expectations on the presence and properties of magnetic fields. The relation between rotation and activity, hence presumably also between rotation and magnetic flux, and the difficulty to detect Zeeman signatures in rotationally broadened spectral lines causes great practical difficulty, especially in sun-like stars. In low-mass (M-type) and pre-main sequence stars, the relation between activity and rotation is presumably more observer-friendly, facilitating the detectability of Zeeman broadening. I will, therefore, distinguish between magnetic field observations in sun-like stars, low-mass stars, and pre-main sequence stars.
The general difficulties detecting the subtle effects of Zeeman broadening in a spectral line from the
spatially unresolved stellar disk were discussed in Section 2.1. A promising way to overcome the problem of
degeneracies between Zeeman broadening and other broadening agents is to compare spectral lines with
different Zeeman sensitivities in the same spectrum. An enhanced width (or equivalent width) of the
magnetically sensitive lines often is good indication for the presence of a magnetic field. This strategy was
successfully applied to Ap stars with fields of 1 kG-strength by Preston (1971). Vogt (1980) used a
multichannel photoelectric Zeeman analyzer mainly to measure polarization in sun-like stars, but also
presents comparison between the widths (FWHM) of magnetically sensitive lines at 6173 Å and two
nearby, magnetically less sensitive lines. Four stars were analyzed with this method finding no
evidence for magnetic fields. Vogt (1980) concludes that this rules out the presence of non-coherent
longitudinal fields in excess of 1000 – 1500 G and covering the entire surface, which is similar to
G.
Robinson Jr (1980) introduced a new method based on the comparison between magnetically sensitive
and insensitive lines. Realizing that the increase in line width for fields less than several kG is very small at
optical wavelengths, he suggested to employ a Fourier transform technique to easily separate the
broadening effects due to magnetism from other broadening effects. The underlying principle
is the very same as if one is comparing line shapes or line widths directly in the wavelength
regime (instead of Fourier regime). However, the Fourier transform technique is able to cleanly
separate the different broadening effects at least in principle, and thus could ideally separate
magnetic broadening from other effects. The main limitation of Zeeman broadening measurements
at optical wavelengths, however, cannot be overcome by this method: it is still necessary to
precisely measure a magnetically non-broadened line in order to use it as a template for the
(potentially stronger) broadening observed in a magnetically sensitive line. Both lines must be of
very similar nature in terms of formation height and temperature response. It is, therefore, not
surprising that the limitations discussed by Robinson Jr (1980) are essentially identical to the
limitations arising when line widths are compared directly. Consequently, the Fourier technique was
not applied to a great many spectra, but the paper became a benchmark for line comparison
techniques in general because it thoroughly discusses the requirements and limitations of this
technique.
What followed was a series of attempts trying to measure magnetic fields in more or less active sun-like
stars. Driven by detections of chromospheric and coronal activity, active stars with relatively low rotational
broadening (
) were observed in order to search for the effects of Zeeman broadening. Highest
obtainable data quality at this time was typically on the order of 
50 000 – 70 000 and
SNR 
100 – 200. A remarkable conclusion from the magnetic field observations taken during this time
was pointed out by Gray (1985). Investigating the reports on magnetic field measurements, he finds that for
G- and K-dwarfs, the product between the magnetic field strength B, and the areal coverage factor f , i.e.,
the average magnetic field strength Bf , “is a constant independent of physical parameters such as
spectral type and rotational velocity”. Realizing that this is rather unlikely, he concludes that
“either we have systematic misconceptions involved in our Zeeman-broadening analysis or else we
have before us a remarkable magnetic conservation condition”. The value of this “magnetic
constant” is roughly Bf = 500 G. According to Equation (4
), this means an extra-broadening of
700 m s–1 for a magnetically sensitive line (
) over an insensitive line (
)
at red optical wavelengths (670 nm); this is typically between 10% and 20% of a resolution
element.
This example demonstrates that searching for the subtle effects of a several hundred Gauss magnetic
field is close to the theoretical detectability of the Zeeman effect, and that it is extremely difficult to judge
whether differences between lines of different magnetic sensitivities are really due to magnetism.
Consequently, the Zeeman analysis methods were criticized by many authors (see e.g., Saar, 1988) centering
on two flaws: 1) incomplete treatment of radiative transfer, and 2) lack of correction for line blends. Saar
(1988) presents a set of improved methods for the analysis of magnetic fields in cool stars. Main ingredients
are radiative transfer effects, treatment of exact Zeeman patterns, and improved correction for line
blends. Following up on this improvement, Basri et al. (1990) went one step further introducing a
two-component analysis by applying their more detailed line-transfer analysis to the (more realistic)
situation in which the magnetic component of the stellar atmosphere is not identical to the
non-magnetic component. The authors also point out that the derived magnetic flux still could be
in error by a factor of 2 because atmospheres from one-dimensional calculations are used for
a multi-component analysis (neglecting gradients and differences in atmospheric structure);
misestimates of abundance, turbulence, and subsequently magnetic field can be quite severe. A detailed
parameter study estimating the accuracy of magnetic field analysis methods in detailed radiative
transfer calculations with embedded fluxtubes is given by Saar and Solanki (1992) and Saar et al.
(1994a).
Obviously, a straightforward way to improve magnetic field measurements is to observe at longer
wavelengths (see Equation (4)). Useful lines are found for example at 1.56 µm (Fe i) and 2.22 µm (Ti i),
i.e., at wavelengths a factor of 3 – 4 longer than typical red/optical observations. First suitable
instrumentation at such long wavelengths became available in the early-1990s. The first detailed
analysis of a high-resolution infrared spectrum in a sun-like star (for earlier work on M stars, see
Section 3.1.2) was performed by Valenti et al. (1995). These authors used a high-resolution
(R = 103 000), high SNR (100 – 200) spectrum (taken during several hours of exposure) to
determine the magnetic field of 
Eri, and upper limits on the order of 100 G in two other
early K-dwarfs. 
Eri has been subject to magnetic field investigations many times earlier at
optical wavelengths. Valenti et al. (1995) also show a compilation of reports on magnetic field
measurements in this star published between 1984 and their work in 1995. Interestingly, average
magnetic fields of 
Eri decreased over time starting at 
800 G in 1984 and reaching 130 G
in 1995. Possible interpretations of this result are that the field in 
Eri is variable, or that
observations reporting lower field strengths (predominantly near-IR measurements) probe a
different part of the stellar atmosphere. Valenti et al. (1995) discuss possible scenarios reaching
the conclusion that probably optical investigations have overestimated the magnetic flux of
Eri.
A critical compilation of magnetic field measurements obtained between the paper of Robinson Jr
(1980) and 1996 was attempted by Saar (1996b). The selection process leading to a condensed sample of
“improved” field measurements was described as follows: “I have therefore compiled a carefully selected
sample of magnetic measurements from analyses which treat radiative transfer effects and use
disk-integration in their models. In addition, I (ruthlessly!) neglect results from low S/N IR data,
measurements using Fe I 8468 Å in K dwarfs, Zeeman/magnetic Doppler imaging results, and
curve-of-growth analyses” (for the reasons why some techniques were neglected, see Saar, 1996b). A similar,
upgraded collection of Zeeman analyses carried out in the period 1996 – 2001 was given by Saar
(2001).
For this review, I have tried in Table 1 to compile magnetic field measurements available for sun-like
stars. Following Saar (1996b), I include only those measurements that rely on relatively high
data-quality and analysis techniques. Since apparently not very many magnetic field measurements
were reported in sun-like stars after 2001, Table 1 does not contain many results in addition
to the compilations by Saar (1996b, 2001). However, in the light of the results reported by
Valenti et al. (1995), I distinguish between work done at optical wavelengths and work done at
infrared wavelengths, the former probably being more prone to overestimating the magnetic
field.
. Green lines indicate the difference
between overall best-fit and the comparison model, i.e., the change in line shape due the presence of
magnetic flux (other fit parameters vary freely). Top: Model with identical temperature for magnetic
and non-magnetic regions; best-fit: Bf = 500 G (solid blue), comparison: Bf = 0 G (dash-dotted
red). Bottom: Best fit for model with different temperatures for magnetic and non-magnetic regions;
Bf = 120 G (blue solid), comparison: solution from upper panel (Bf = 500 G, same temperatures,
red dashed line) (from Anderson et al., 2010, reprinted with permission āø ESO). A critical re-investigation of the detectability of magnetic fields in high-quality optical spectra was
carried out by Anderson et al. (2010). The data material used for this work is of much higher quality than
most magnetic field investigations before, and the data therefore allows a critical view on the published
results and some of the limitations of the method. Anderson et al. (2010) used optical spectra around the
Fe i line at 6173 Å observed at a spectral resolving power of R = 220 000 and SNR 
400. The
analysis is carried out for a one-component model with the same atmosphere for the magnetic and
the non-magnetic parts of the stellar surface, and also for a two-component model employing
different atmospheres for the two components. The results are reproduced in Figures 12
and 13.
For the active G0 star 59 Vir, the authors find a magnetic field with Bf 
420 G for the
one-component case. For the two-component analysis, they cannot exclude a zero-field solution
reporting an upper limit of 300 G. Figure 12 shows how subtle the differences between solutions
with different magnetic field strengths are if all other relevant parameters are allowed to vary
freely (there is currently no way to constrain these parameters at the level required). Figure 13
demonstrates the relation between magnetic field strength B and filling factor f in case of a
one-component atmosphere (left panel). The two-component models shown in the center and right panels,
however, can lift the Bf degeneracy but manage to reproduce the spectra even without the
presence of a significant magnetic field. In other words, at optical wavelengths, the signal of
temperature spots on the surface of a cool star can dominate the influence of the magnetic field
through Zeeman broadening. Unfortunately, we have so far no clear empirical evidence for the
relation between temperature and magnetic field strength on stellar surfaces other than on the
Sun.
A look at Table 1 reveals that infrared measurements are only available in six sun-likes stars, all of them are of spectral type K. Two of the six data points are actually non-detections, and three were reported in conference summaries in which, unfortunately, no comprehensive presentation of the data and its analysis is given.
-maps for 59 Vir solutions (see Figure 12). Left panel: Solution using the same
atmospheres for magnetic and non-magnetic regions; Center panel: the same using cool magnetic
regions; Right panel: the same for warm magnetic regions (from Anderson et al., 2010, reprinted with
permission āø ESO).Low-mass stars of spectral type M have radii of approximately half a solar radius and less. If the
stellar dynamo depends on the value of the Rossby number, 
, the magnetic field
strength expected in sun-like and low-mass stars is a function of rotational period and convective
overturn time. Values for the convective overturn time are theoretically not well determined, but
is probably higher at lower masses (e.g., Kim and Demarque, 1996). Therefore, slower
rotation is sufficient to produce larger fields in less massive stars. Furthermore, the smaller radii of
less massive stars lead to lower surface velocities hence less rotational broadening at a given
rotational period. Finally, less massive stars are also much cooler thus exhibiting less temperature
broadening in their spectral lines. It is this combination of parameters that facilitates the detection
of Zeeman splitting in M-type stars in comparison to more massive, sun-like stars; Zeeman
broadening is more easily detected because of generally narrower line widths (see also Reiners,
2007).
). The first detection of Zeeman splitting in an M-type star, and also the first detection of a photospheric
magnetic field in cool stars at all, was presented by Saar and Linsky (1985). They observed the early-M flare
star AD Leo using a Fourier transform spectrometer. After six hours of observation they had obtained a
spectrum with R = 45 000 and SNR 
25 around the Ti i lines at 2.22 µm from which they measured
an average magnetic field strength of Bf = 2800 G. Similar data taken with the same instrument
was obtained in a few M-stars, and Saar (1994) presented a preliminary analysis of the three
M-type stars AU Mic, AD Leo, and EV Lac. Another benchmark was the investigation of
the Fe i line at 8468 Å in seven early- to mid-M dwarfs by Johns-Krull and Valenti (1996).
Substantial magnetic fields were detected in two stars of the sample, EV Lac and Gl 729. A refined
analysis of the two stars and AD Leo and YZ Cmi was presented in Johns-Krull and Valenti
(2000). The latter work assumed a distribution of magnetic fields on the stellar surface, which led
to significantly higher average field values compared to Johns-Krull and Valenti (1996). The
results from the 8468 Å line were comparable to the values from the 2.22 µm line within
10 – 20%.
A serious problem for the detection of Zeeman splitting in atomic spectral lines of M-type stars is the
appearance of molecular bands. For example, the Fe i line at 8468 Å is embedded in a forest of TiO
molecular absorption lines, which makes the modeling of Zeeman splitting in this line a delicate task. To
overcome this problem, and the notorious difficulty to model TiO absorption (see Valenti et al.,
1998), Johns-Krull and Valenti (1996) modeled the ratio of the flux between an active and an
inactive star. Hopefully, our understanding of very cool atmospheres, molecular chemistry, and
molecular line formation will in the future allow a detailed modeling of Zeeman splitting in the
spectra of M dwarfs (see also Kochukhov et al., 2009; Önehag et al., 2011; Shulyak et al.,
2011).
At optical wavelengths, the main opacity contributors in M dwarfs are molecular bands from TiO and
VO. Analysis of Zeeman broadening in these bands, however, is difficult not only because of problems
getting the line formation right, but also because the lines are not individually resolved. Nevertheless, for
the detection of M star magnetic fields, it would be favorable to utilize molecular absorption bands. A
molecular band that appears to be extremely useful for the analysis of M-star magnetic fields (and other
purpose) is the near-infrared band of molecular FeH. Its suitability for magnetic analysis was shown by
Wallace et al. (1999), and it was proposed to be a useful diagnostic at low temperatures by Valenti
et al. (2001). An observational problem of FeH is that its most suitable band is located at
around 
= 1µm, which is too red for most CCDs and too blue for most astronomically
used infrared spectrographs. As a consequence, only very few high-resolution spectrographs can
provide spectra at this wavelength, and efficiencies are typically ridiculously low. On the other
hand, M dwarfs emit much of their flux at near-infrared wavelengths so that in comparison to
optical measurements, the signal quality around 1µm is not much lower than around 700 nm
if the spectra are obtained with an optical/near-IR echelle spectrograph like HIRES (Keck
observatory) or UVES (ESO/VLT). Reiners and Basri (2006) developed a method to semi-empirically
determine the magnetic fields of M dwarfs comparing FeH spectra of the targets to spectra of
two template stars; one with no magnetic field and one with a known, strong magnetic field
(Figure 15). This method requires a known magnetic star to calibrate the Zeeman splitting amplitude.
The field strength of the target star is then estimated by interpolation between the template
spectra.
).
The black histogram shows the spectrum of Gl 729. The red and blue lines are scaled spectra of the
active star EV Lac and the inactive star GJ 1002, respectively. The green line is an interpolation
between the red (Bf = 3.9 kG) and blue (0 kG) lines yielding a field strength of Bf = 2.2 kG for
Gl 729 (from Reiners and Basri, 2007, reproduced by permission of the AAS). The method of Reiners and Basri (2006) was first used in a sample of 24 M stars between spectral types
M0 and M9 (Reiners and Basri, 2007). As reference, the field measurement of EV Lac measured by
Johns-Krull and Valenti (2000) was used. Thus, all magnetic field measurements are relative to this
reference star (
= 3.9 kG), and magnetic fields higher than this value cannot be quantified.
Obviously, systematic uncertainties of the measurements are quite large, typically several hundred Gauss,
and uncertainties probably grow towards very late spectral types where the template spectra are less suited
as a reference. Unfortunately, Zeeman splitting of the FeH molecule is very complicated and could not
entirely be described at this point (see Berdyugina and Solanki, 2002). Meanwhile, progress has been made
using an empirical approach to understand FeH absorption and line formation (Wende et al., 2009,
2010), and to model Zeeman splitting in FeH lines (Afram et al., 2009; Shulyak et al., 2010).
It was suggested that the fields determined semi-empirically may be overestimated by some
20%1
(Shulyak et al., 2010).
A (probably non-exhaustive) list of magnetic field measurements from Stokes I analysis in M dwarfs is
given in Table 2, and I plot the distribution of field strength as a function of spectral type in
Figure 16. The field strengths of young, early-M and field mid- and late-M dwarfs are on the
order of a few kG. This is the main results from Zeeman analysis and consistently found using
different indicators (at least in mid-M dwarfs). Compared to the Sun, the average magnetic field
hence is larger by two to three orders of magnitude, an observational result that must have
severe implications for our understanding of low-mass stellar activity. It is not clear whether
our picture of a star with spots more or less distributed over the stellar surface is actually
valid in M dwarfs. If, for example, 50% of the surface of a star with a mean magnetic field of
4 kG is covered with a “quiet” photosphere and low magnetic field, the other half of the star
must have a field strength as large as 
8 kG. The two components of the stellar surface
on such a star probably have very different temperatures and properties, and the definition
of effective temperature must be considerably different from the temperature of the “quiet”
photosphere.
In early-M dwarfs (M3 and earlier), magnetic fields were found in young stars that are still rapidly rotating. Since old, early-M dwarfs in the field are generally slowly rotating and inactive there has been no search for magnetic fields in any large sample of them. Typical field values can be expected to be on the level of a few hundred Gauss and less, which is difficult to detect with Stokes I Zeeman measurements. Many mid- and late-M stars are rapidly rotating and fields of kG-strength are ubiquitously found among them.
Magnetic fields of pre-main sequence stars are of particular interest because accretion of circumstellar
material onto the stellar surface is believed to be controlled by the stellar magnetic field (e.g., Bouvier
et al., 2007). Evidence for accretion is observed in pre-main sequence stars of very different mass including
young brown dwarfs. Field strengths predicted from several models of magnetospheric accretion are on
the order of several kG for T Tauri stars, and a few hundred Gauss for young brown dwarfs
(Johns-Krull et al., 1999b; Reiners et al., 2009b). On top of this, at young ages, magnetic fields
may be generated by a dynamo like in older, sun-like stars (in contrast to fossil fields), but
the dynamo would probably operate similar to the one in low-mass M-type dwarfs because
pre-main sequence stars are still fully convective. On the other hand, at ages of a few Myr,
primordial fields may still be present and not (yet) dissipated. Magnetic fields in pre-main
sequence stars may therefore carry important information about the star- and planet-formation
process.
First measurements of magnetic field strengths in T Tauri stars were attempted using the equivalent
width method in (red) optical absorption lines by Basri et al. (1992), and Guenther et al. (1999) were
following this strategy. Johns-Krull et al. (1999b) used infrared lines of Ti i at 2.22 µm to determine the
magnetic field in BP Tau. Obtaining information on stellar parameters and rotation from optical lines and
magnetically insensitive CO lines, they were able to disentangle the significant Zeeman broadening from
other broadening agents. Similar work on other T Tauri stars using infrared spectra was done by
Johns-Krull et al., much of it is summarized in Johns-Krull (2007) where additional measurements of 14
T Tauri star magnetic fields are presented. Another set of 14 magnetic field measurements in very young
T Tauri stars in the Orion nebula cluster are given in Yang and Johns-Krull (2011). We will return to the
results from these substantial samples in Sections 5 and 7.3. A summary of magnetic field
measurements in very young stars and brown dwarfs is given in Table 3, and are shown in
Figure 17.
Using FeH measurements as laid out in Section 3.1.2, Reiners et al. (2009b) attempted to find evidence
for kG-strength magnetic fields in young brown dwarfs. Young brown dwarfs can be expected to harbor
substantial magnetic fields (Reiners and Christensen, 2010), and no fundamental difference is known to exist
in the parameters that are believed to be relevant for magnetic flux generation between very-low mass stars
and young brown dwarfs. However, in contrast to pre-main sequence stars, and in contrast to older brown
dwarfs, none of the young brown dwarfs investigated by Reiners et al. (2009b) exhibited a field above
the detection threshold that in all cases lay below the fields typically found among the other
groups.
An important property of the four young brown dwarfs investigated for magnetic fields is that all of them show evidence for accretion and, therefore, harbor a circumstellar disk. Magnetic field strengths required for magnetospheric accretion in these objects are much lower than in more massive, young stars, hence there is currently no contradiction between the presence of accretion and the lack of evidence for substantial fields. Observations of radio-emission, however, indicate that fields of a few kG strength are in fact present on some L-type field (old and non-accreting) brown dwarfs (Hallinan et al., 2008; Berger et al., 2009). Direct measurement of magnetism in non-accreting brown dwarfs, both young and old, are required to further investigate whether the average fields are really weaker in young brown dwarfs or in the presence of accretion.
It is an interesting question whether the non-detection of magnetic fields in brown dwarfs is due to the presence of accretion disks around the objects observed so far. If this is the case, there ought to be some mechanism for the disk to regulate the magnetic field of the central object, which is not easily understood. Alternatively, large difference in radius may be of importance in this context because the surface area of young brown dwarfs is about an order of magnitude larger than the surface of old brown dwarfs. If magnetic flux is approximately conserved during its evolution, the average magnetic field would be an order of magnitude lower in young, large brown dwarfs than in old, small, field brown dwarfs. We come back to this point in Section 7.3.
| Star | Other Name | SpType | Bf [kG] | Reference |
| Gl 182 | M0.0 | 2.5 | Reiners and Basri (2009) |
|
| Gl 803 | AU Mic | M1.0 | 2.3 | Saar (1994) |
| Gl 569A | M2.0 | 1.8 | Reiners and Basri (2009) |
|
| Gl 494 | DT Vir | M2.0 | 1.5 | Saar (1996b) |
| Gl 70 | M2.0 | < 0.2 | Reiners and Basri (2007) |
|
| Gl 873 | EV Lac | M3.5 | 3.9 | Johns-Krull and Valenti (1996, 2000) |
| 3.7 | Saar (1994) |
|||
| Gl 729 | M3.5 | 2.0 | Johns-Krull and Valenti (1996, 2000) |
|
| 2.2 | Reiners and Basri (2007) |
|||
| Gl 87 | M3.5 | 3.9 | Reiners and Basri (2007) |
|
| Gl 388 | AD Leo | M3.5 | 2.8 | Saar and Linsky (1985) |
| 2.6 | Saar (1994) | |||
| 3.3 | Johns-Krull and Valenti (2000) |
|||
| 2.9 | Reiners and Basri (2007) |
|||
| 3.2 | Kochukhov et al. (2009) |
|||
| GJ 3379 | M3.5 | 2.3 | Reiners et al. (2009a) |
|
| GJ 2069 B | M4.0 | 2.7 | Reiners et al. (2009a) |
|
| Gl 876 | M4.0 | < 0.2 | Reiners and Basri (2007) |
|
| GJ 1005A | M4.0 | < 0.2 | Reiners and Basri (2007) |
|
| Gl 490 B | G 164-31 | M4.0 | 3.2 | Phan-Bao et al. (2009) |
| Gl 493.1 | M4.5 | 2.1 | Reiners et al. (2009a) |
|
| GJ 4053 | LHS 3376 | M4.5 | 2.0 | Reiners et al. (2009a) |
| GJ 299 | M4.5 | 0.5 | Reiners and Basri (2007) |
|
| GJ 1227 | M4.5 | < 0.2 | Reiners and Basri (2007) |
|
| GJ 1224 | M4.5 | 2.7 | Reiners and Basri (2007) |
|
| Gl 285 | YZ CMi | M4.5 | 3.3 | Johns-Krull and Valenti (2000) |
| > 3.9 | Reiners and Basri (2007) |
|||
| 4.5 | Kochukhov et al. (2009) | |||
| GJ 1154 A | M5.0 | 2.1 | Reiners et al. (2009a) |
|
| GJ 1156 | M5.0 | 2.1 | Reiners et al. (2009a) |
|
| Gl 905 | M5.5 | < 0.1 | Reiners and Basri (2007) |
|
| GJ 1057 | M5.0 | < 0.2 | Reiners and Basri (2007) |
|
| Gl 905 | M5.5 | < 0.1 | Reiners and Basri (2007) |
|
| GJ 1245B | M5.5 | 1.7 | Reiners and Basri (2007) |
|
| GJ 1286 | M5.5 | 0.4 | Reiners and Basri (2007) |
|
| GJ 1002 | M5.5 | < 0.2 | Reiners and Basri (2007) |
|
| Gl 406 | M5.5 | 2.4 | Reiners and Basri (2007) |
|
| 2.1 – 2.4 | Reiners et al. (2007) | |||
| Gl 412 B | M6.0 | > 3.9 | Reiners et al. (2009a) |
|
| GJ 1111 | M6.0 | 1.7 | Reiners and Basri (2007) |
|
| Gl 644 C | VB 8 | M7.0 | 2.3 | Reiners and Basri (2007) |
| GJ 3877 | LHS 3003 | M7.0 | 1.5 | Reiners and Basri (2007) |
| 2M 0440232–053008 | M7.0 | 1.6 | Reiners and Basri (2010) |
|
| 2M 0741068+173845 | M7.0 | 1.0 | Reiners and Basri (2010) |
|
| 2M 0752239+161215 | M7.0 | 3.5 | Reiners and Basri (2010) |
|
| 2M 0818580+233352 | M7.0 | 1.0 | Reiners and Basri (2010) |
|
| 2M 1048126–112009 | GJ 3622 | M7.0 | 0.6 | Reiners and Basri (2010) |
| 2M 1356414+434258 | M7.0 | 2.7 | Reiners and Basri (2010) |
|
| 2M 1456383–280947 | M7.0 | 1.2 | Reiners and Basri (2010) |
|
| 2M 1534570–141848 | M7.0 | 2.0 | Reiners and Basri (2010) |
|
| LHS 2645 | M7.5 | 2.1 | Reiners and Basri (2007) |
|
| 2M 0331302–304238 | M7.5 | 2.0 | Reiners and Basri (2010) |
|
| 2M 0351000–005244 | M7.5 | 1.4 | Reiners and Basri (2010) |
|
| 2M 0417374–080000 | M7.5 | 1.8 | Reiners and Basri (2010) |
|
| 2M 0429184–312356A | M7.5 | 2.5 | Reiners and Basri (2010) |
|
| 2M 1006319–165326 | M7.5 | 1.6 | Reiners and Basri (2010) |
|
| 2M 1246517+314811 | M7.5 | < 0.4 | Reiners and Basri (2010) |
|
| 2M 1253124+403403 | M7.5 | 1.6 | Reiners and Basri (2010) |
|
| 2M 1332244–044112 | M7.5 | 1.6 | Reiners and Basri (2010) |
|
| 2M 1546054+374946 | M7.5 | 2.7 | Reiners and Basri (2010) |
|
| LP 412-31 | M8.0 | > 3.9 | Reiners and Basri (2007) |
|
| VB 10 | M8.0 | 1.3 | Reiners and Basri (2007) |
|
| 2M 0248410–165121 | M8.0 | 1.4 | Reiners and Basri (2010) |
|
| 2M 0320596+185423 | M8.0 | 3.7 | Reiners and Basri (2010) |
|
| 2M 0517376–334902 | M8.0 | 1.6 | Reiners and Basri (2010) |
|
| 2M 0544115–243301 | M8.0 | 1.2 | Reiners and Basri (2010) |
|
| 2M 1016347+275149 | M8.0 | 2.1 | Reiners and Basri (2010) |
|
| 2M 1024099+181553 | M8.0 | < 1.4 | Reiners and Basri (2010) |
|
| 2M 1141440–223215 | M8.0 | 1.8 | Reiners and Basri (2010) |
|
| 2M 1309218–233035 | M8.0 | 1.2 | Reiners and Basri (2010) |
|
| 2M 1440229+133923 | M8.0 | < 0.6 | Reiners and Basri (2010) |
|
| 2M 1843221+404021 | M8.0 | 1.2 | Reiners and Basri (2010) |
|
| 2M 2037071–113756 | M8.0 | < 0.2 | Reiners and Basri (2010) |
|
| 2M 2306292–050227 | M8.0 | 0.6 | Reiners and Basri (2010) |
|
| 2M 2349489+122438 | M8.0 | 1.2 | Reiners and Basri (2010) |
|
| 2M 0024442–270825B | M8.5 | 2.1 | Reiners and Basri (2010) |
|
| 2M 0306115–364753 | M8.5 | 1.6 | Reiners and Basri (2010) |
|
| 2M 1124048+380805 | M8.5 | 2.0 | Reiners and Basri (2010) |
|
| 2M 1403223+300754 | M8.5 | 2.1 | Reiners and Basri (2010) |
|
| 2M 2226443–750342 | M8.5 | 1.8 | Reiners and Basri (2010) |
|
| 2M 2331217–274949 | M8.5 | 3.1 | Reiners and Basri (2010) |
|
| 2M 2353594–083331 | M8.5 | 2.0 | Reiners and Basri (2010) |
|
| LHS 2924 | M9.0 | 1.6 | Reiners and Basri (2007) |
|
| LHS 2065 | M9.0 | > 3.9 | Reiners and Basri (2007) |
|
| 2M 0019457+521317 | M9.0 | 3.7 | Reiners and Basri (2010) |
|
| 2M 0109511–034326 | M9.0 | 1.4 | Reiners and Basri (2010) |
|
| 2M 0334114–495334 | M9.0 | 1.4 | Reiners and Basri (2010) |
|
| 2M 0443376+000205 | M9.0 | < 1.0 | Reiners and Basri (2010) |
|
| 2M 0853362–032932 | M9.0 | 2.9 | Reiners and Basri (2010) |
|
| 2M 1048147–395606 | M9.0 | 2.3 | Reiners and Basri (2010) |
|
| 2M 1224522–123835 | M9.0 | 1.4 | Reiners and Basri (2010) |
|
| 2M 1438082+640836 | M9.5 | 1.2 | Reiners and Basri (2010) |
|
| 2M 2237325+392239 | M9.5 | 1.0 | Reiners and Basri (2010) |
|
| Publication | Comment |
| Robinson et al. (1980) |  Boo A, 70 Oph A, 61 Vir |
| Vogt (1980) | BY Dra, HD 88230, 61 Cyg A, HD 209813 |
| Marcy (1981) |  Boo A |
| Golub et al. (1983) |  And |
| Marcy (1984) | 29 G- and K main sequence stars |
| Marcy and Bruning (1984) | 8 evolved stars |
| Gray (1984) | 18 F-, G-, and K-dwarfs |
| Gondoin et al. (1985) |  Boo A, 61 UMa,  And |
| Saar et al. (1986) | EQ Vir |
| Bruning et al. (1987) | 7 K- and M-dwarfs |
| Mathys and Solanki (1989) | Preliminary results, “Stenflo–Lindegren” technique |
| Bopp et al. (1989) | VY Ari |
| Saar (1990) |
Selection of 31 G – M star measurements; including unpublished data |
 |
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|
Living Rev. Solar Phys. 8, (2012), 1
http://www.livingreviews.org/lrsp-2012-1 |
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