Download Mensajero Estelar

Survey
yes no Was this document useful for you?
   Thank you for your participation!

* Your assessment is very important for improving the work of artificial intelligence, which forms the content of this project

Document related concepts
no text concepts found
Transcript
SOVAFA
ACA
Sociedad Venezolana
de Aficionados a la
Astronomía
Asociación Carabobeña de
Astronomía
Mensajero
Estelar
Año 39
Nº 72
Octubre- Diciembre de 2014
Contenido:
-
Noticias
- ¿Llego el Voyager a la Heliopausa?
Radiantes del Trimestre
- La Supernova más brillante de la historia
Fases de la Luna
- Infrared Photometry of the Pleyades
Cúmulo Estelar de las Híades
- Nuevo Ciclo del Calendario Maya
Eclipse Total de Luna de Oct. 08
- ¿Por qué la Luna no es una esfera perfecta?
¿Cuánto de la superficie lunar vemos?
- Methane Plumes in the Arctic
Solar Variability and Terrestral Climate
- Olas de 5 m de altura erosionan el Hielo Ártico
¿Pudo el Bosón de Higgs Colapsar el U.
- Old pre-main-sequence stars
Ocultación de Marte por la Luna
- Temperaturas Anómalas en Julio y Agosto
La GMR de Júpiter se achica
- Deflexión de la luz por el Sol
Nuevas Enanas Rojas cercanas al Sol
- Born Betwen Nov. 29 and Dec. 18 …
A new wiew of the red planet
- Planetas Acuosos
Violenta historia del Sol joven…
- Meteorito en Nicaragua
Junio de 2014, el más cálido registrado
- Técnica Lucky Image…
- Las Geminíadas
www.sovafa.com, www.sovafa.org, [email protected], @astrorecord, @sovafa
Noticias
1.- La Corona Solar es mucho más grande que lo que se pensaba. Recientes observaciones realizadas con el satélite
STEREO evidencian que esta se extiende más de 8 millones de km de la superficie del Sol.
2.- Datos obtenidos por la sonda Cassini parecen indicar que el interior de Titán podría contener un océano muy salado,
de acuerdo a mediciones de micro gravedad realizadas por la sonda. Estas sales serían de Azufre, Sodio, y Potasio.
3.- El Observatorio “Athena” fue incorporado a la visión cósmica de la Agencia Espacial Europea, ESA, en su plan 2015
– 2025. El mismo estudiará el Universo caliente y energético, estará en el punto de equilibrio gravitatorio “Lagrange 2
(L2), previéndose su lanzamiento para el año 2028.
4.- Un equipo de investigadores de la Universidad de Nueva Gales del Sur, en Australia, ha descubierto un planeta similar
a la Tierra potencialmente habitable a tan sólo 16 años luz de distancia. Llamado Gliese 832, es una "súper-Tierra" con
una masa 5 veces la de nuestro planeta. Tarda 16 días en completar una órbita alrededor de su estrella; una enana roja
cuyo brillo es menor al del Sol, pero debido al tiempo que tarda en orbitarla tiene aproximadamente la misma energía
estelar que la Tierra.
5.- En Julio la actividad solar volvió a disminuir muy drásticamente. Luego del Máximo a que llegó el año pasado, esta
actividad decayó de manera bastante brusca, para luego, desde abril pasado volver a incrementarse, y en Julio volvió a
caer de manera brusca.
6.- El asteroide 2014 KM4 de 192 metros descubierto, a principios mayo, ha pasado de forma segura por el sistema
Tierra-Luna a 0.17 AU de distancia el 21 de abril. Hasta el momento, la trayectoria lo lleva a una ruta de colisión con el
gigante del Sistema Solar, Júpiter en el año 2022.
7.- El Exoplaneta OGLE-2013-BLG-0341LBb, situado a 3.000 A.L. de la Tierra, posee condiciones muy parecidas a las
del entorno terrestre. Es más frío que la Tierra, pues su estrella es más pequeña que el Sol, pero es un objeto
potencialmente habitable.
8.- El 2 de Septiembre se descubrió un asteroide que nos pasó a unos 38.000 km dos días antes. El objeto mide solo unos
8 m de diámetro y fue catalogado como 2014 RA
9.- El día 3 de Septiembre se observaron 305 bolas de fuego sobre el SE de USA, el número más elevado que se ha
observado hasta ahora. Este año solo el día 13 de Agosto, máximo de las Perseidas se detectaron unos 105 bólidos, y la
gran mayoría de ellos conectados con el radiante.
10.- En Agosto la periodista y locutora Amalia Heller entrevistó a Jesús Otero por Mágica 99.1 FM sobre la lluvia de
estrellas de las Perseidas y la Conjunción de Venus y Júpiter, en su programa La Magia de Amalia Heller que se transmite
de Lunes a Viernes de 7:00 a 8:30 pm
11.- El 04 de Septiembre Amalia Heller nuevamente entrevistó a Jesús Otero, pero esta vez sobre el Asteroide 2014 RC
que nos pasó rozando el día 07 de septiembre.
12.- El 31 de Agosto y el 07 de Septiembre 2 asteroides pasaron a unos 40.000 Km de la Tierra, esto es unos 33.000 km
de la Superficie terrestre, sabemos que estos pasos rasantes son comunes, pero también sabemos que en algún momento
seremos impactados.
13.- Un instrumento de la NASA abordo del orbitador Rosetta de la ESA, ALICE, descubrió que el cometa
67P/Churyumov-Gerasimenko es inusualmente oscuro, muy oscuro. Analizando las ondas ultravioletas de la superficie
del cometa, también detectó la presencia de hidrogeno y oxigeno en su coma y pocas evidencias de hielo de agua
expuesto.
14.- El 6 de septiembre cayó un meteorito en Nicaragua. Hubo un fuerte destello en el cielo, luego una fuerte explosión, e
instantes después un sismo suave producido por el impacto. El objeto dejó un cráter de 12 m de diámetro y unos 5 de
profundidad, el objeto pudo haber tenido casi un metro de diámetro y debe encontrarse bajo los rellenos post impacto del
cráter.
15.- El día 06 de Septiembre se observó un bólido de magnitud -12 con ruido sobre la ciudad de Barcelona, España. En
Agosto y Septiembre la Tierra estuvo pasando por zonas muy densas de restos interplanetarios. De hecho el 03 de
Septiembre se contabilizaron 305 Bolas de Fuego, solo en el SE de USA, y el 06 hubo 68.
16.- El 01 de Sepiembre, en Campinas, Brazil 4 brillantes bolas de fuego cruzaron el cielo con magnitude de -4 a -8,
Agosto y Septiembre registraron muchos meteoros brillantes y la caída de un objeto al Norte de Managua, Nicaragua.
17.- Estudios recientes del USGS demuestran que es alktamente improbable que el volcán de Yosemite haga erpción en el
futuro cercano.
18.- Evidencia de tectónica de placas en la Luna Europa de Júpiter. Se observó lo que parece una tectónica de Subducción
en las capas de hielo del satélite.
19.- La Vía Láctea forma parte de un supercúmulo que ha sido bautizado como Supercúmulo de Laniakea, o cielo
inmenso, que tiene 500 millones de Años Luz de Grosor. Esto fue descubierto con el Green Bank Radiotelescope. Este
filamento contiene al menos 100.000 galaxias.
Radiantes del Trimestre
Radiante
Oriónidas
Taúridas del Sur
Taúridas del N.
Leónidas
Androménidas
5185C.Minóridas
51 Androménidas
Piscidas
43 Taúridas
Geminíadas
Púpidas - Vélidas
Fecha
Octubre 17 - 26
Sept.15-Nov 30
Sept. 19-Dic. 5
Nov. 14 - 20
Nov. 4 - 20
Dic. 1 - 5
Diciembre 04
Dic. 10 - 14
Dic. ¿? - 13
Dic. 13 - 16
Nov. 24 - Ene 9
Máximo
Oct. 19 - 23
Nov. 3
Nov. 13
Nov. 17 - 18
Nov. 16
Dic. 3 - 4
Dic. 04
Dic. 10
Dic. 11
Dic. 12 - 13
Dic. 25
T. H. Z.
20
7
9
Var.
Var.
Var.
¿40?
+80
¿97?
145
15
A. R.
06h 18m
03h 22m
03h 53m
10h 12m
01h 44m
07h 36m

+ 15º
+ 13.6º
+ 22º
+ 22º
+ 25º
+ 4º
04h 10m
07h 28m
09h 03m
+19.5º
+ 33º
- 48º
Hora
02:00
23:00
23:00
02:00
21:00
22:00
19:30
19:00
20:00
22:30
00:00
Las Leónidas, las 43 Taúridas, y las Geminíadas son radiantes que dan meteoros brillantes y su número puede
variar mucho de un año a otro.
Las lluvias de estrellas aquí listadas se encuentran todas activas, algunas de ellas son de difícil observación pues
sus meteoros son de poco brillo.
Hay que ver cuál es la fase lunar el día de la observación, pues la luz de la Luna puede afectar mucho la
observación del radiante.
Máximo es el día en que se espera que la lluvia de estrellas llegue a su máximo número de meteoros.
THZ es el número de meteoros que veríamos del radiante si este se encontrara en el zenit.
α y δ son Ascensión Recta y Declinación.
Hora se refiere a la hora en la cual puede empezar a observarse el radiante. Viene en Hora Legal de Venezuela.
O -4,5h GMT.
Este año las Oriónidas ocurrirán entre Cuarto Menguante y Luna Llena.
Geminíadas y 43 Taúridas no serán molestadas por la Luna.
Las Taúridas del Sur serán molestadas por Luna casi Llena, al igual que 51 Androménidas.
Las Púpidas - Vélidas ocurrirán en Luna Nueva y la Luna no interferirá en su observación
Si observa cualquiera de estos radiantes o una actividad meteórica inusual envíe un informe a
[email protected] o un mensaje al Twitter: αastrorecord
Fases de la Luna

Luna Nueva
Fecha
Hora
Sept. 24
06:12
Oct. 23
21:55 P
Nov. 22
12:31
Dic. 22
01:35

Cuarto Creciente
Fecha
Hora
Oct. 01 19:32
Oct. 31 02:48
Nov. 29 10:06
Dic. 28 18:32

Luna Llena
Fecha
Hora
Oct. 08
10:49 t
Nov. 06 22:22
Dic. 06
12:26
Ene. 05
04:53

Cuarto Menguante
Fecha Hora
Oct.15 19:12
Nov. 14 15:17
Dic.14 12:53
Ene. 13 09:48
En Luna Nueva la Luna no se puede ver, pues está en Conjunción con el Sol.
En Cuarto Creciente la Luna se observa en la Tarde y primeras horas de la noche.
En Luna Llena la Luna sale al ocultarse el Sol y se observa durante toda la noche.
En Cuarto Menguante la Luna sale tarde, se observa de madrugada y primeras horas de la mañana.
Estos datos son muy importantes a la hora de planificar sus observaciones, ya sean planetarias, de radiantes u
objetos de espacio profundo.
Téngalas en cuenta para la observación de eventos astronómicos.
t = Eclipse Total de Luna y A = Eclipse Anular de Sol
El Eclipse Total de Luna de Octubre 08 podrá observarse alto en el firmamento en el momento de la totalidad.
Este es un proyecto importante de observación y estamos involucrados en un proyecto internacional.
P Significa Eclipse Parcial de Sol. No será visible en Venezuela.
Cúmulo Estelar de las Híades
Con la excepción del
cúmulo de la Osa Mayor, las
Híades es el cúmulo estelar más
cercano a la Tierra, se encuentra a
una distancia de 151 Años Luz
(AL).
Es
muy
fácil
de
identificar en el firmamento por
ser un cúmulo compacto y con una
forma de V muy característica, lo
que lo hace un Asterismo.
El cúmulo se ve alto en
nuestro firmamento, entre las
Pléyades y la Constelación de
Orión. Utilizando la enfilación de
las estrellas del Cinturón de Orión
en sentido Alnitak – Mintaka y
prolongando esta enfilación hacia
el NW llegamos a él.
La estrella más brillante
que vemos en esta V de estrellas es la gigante roja Aldebarán, cuyo nombre
significa el Ojo del Toro, y que en realidad no forma parte de él. Esta estrella
cierra la V al Sur.
De este cúmulo podemos observar a simple vista poco más de una
docena de estrellas, pero varias decenas de ellas pueden observarse con
pequeños binoculares.
Las Híades eran ninfas, según la mitología griega e hijas de Atlas y
Aethra, las cuales lloraban eternamente a su hermano Hyas, quien fue muerto
por un León. Las Híades eran medio hermanas de las Pléyades, hijas de Atlas
con Pleione. Los dioses colocaron a las Híades y a las Pléyades en el
firmamento a propósito, para salvarlas de los deseos lujuriosos de Orión. Al
mismo tiempo convirtió a Hyas en la constelación de Acuario, y al León que
lo mató en la Constelación de Leo, en la parte opuesta del firmamento.
Así cuando una constelación salía al Este, la otra se ocultaba al Oeste.
Este
mito
muestra
una
ambivalencia con el mito de
Orión y el Escorpión Celeste
enviado por la Diosa Diana para
que picara y diera muerte a
Orión.
Zeus da la inmortalidad
a su hijo Orión convirtiéndolo en
Constelación
y
ubica
al
Escorpión en el lado opuesto del
cielo a fin de evitar un nuevo
encuentro. Así cuando Orión sale
al Este, Scorpio desaparece por el
Oeste y viceversa. De esta
manera nunca podemos ver
ambas constelaciones en el
firmamento al mismo tiempo.
Con
un
telescopio
se
observan cerca de 100 estrellas de
este cúmulo estelar.
Si bien en la Mitología las
Híades y las Pléyades son hermanas, en la
realidad ambos cúmulos son muy
diferentes. Las Pléyades poseen estrellas
Azules muy jóvenes y su edad no supera
los 100 millones de años, por su parte,
Las Híades poseen muchas estrellas rojas,
gigantes rojas y naranjas, así como enanas
blancas, por lo que la edad de este cúmulo
estelar es de más de 700 millones de años.
Pero hay un cúmulo estelar
abierto cuyas características son muy
similares al de las Híades, es el Cúmulo
del Pesebre, en Cáncer. Ambos cúmulos
se mueven en dirección idéntica en el
espacio y sus edades son similares.
Algunos astrónomos creen que a pesar de
encontrarse muy lejos uno del otro, ambos
se formaron en la misma nebulosa hace
unos 700 u 800 millones de años.
Eclipse Total de Luna, Oct. 08, 2014
El 08 de Octubre de 2014 ocurrirá un Eclipse Total de Luna que será visible en Venezuela, desdichadamente la
fase de totalidad no será visible desde Venezuela, pues la Luna se ocultará poco antes de la llegada de la Totalidad. Sin
embargo podremos observar como la Luna se irá eclipsando mientras baja en el horizonte.
Tiempos para Venezuela
El Eclipse Penumbral Empieza:
El Eclipse Parcial Empieza:
El Eclipse Total Empieza:
El Medio del Eclipse:
El Eclipse Total Finaliza:
El Eclipse Parcial Finaliza:
El Eclipse Penumbral Termina:
08:15:33 UT
09:14:48 UT
10:25:10 UT
10:54:36 UT
11:24:00 UT
12:34:21 UT
13:33:43 UT
03:45:33 HLV
04:44:48 HLV
05:55:10 HLV
06:24:36 HLV
06:54:00 HLV:
08:04:21 HLV
09:03:43 HLV
Los próximos Eclipses Lunares observables en Venezuela ocurrirán en Abril 04 de 2015, pero solo podrá el
Inicio del Eclipse y eso con suerte, pues a Luna estará muy cerca del Horizonte. El siguiente será el 28 de Septiembre,
este si será visible en su totalidad desde Venezuela.
El Eclipse de Octubre 08 nos sirve para practicar la observación de este fenómeno, el paso de la sombra sobre
cráteres, y mares, y realizar observaciones. El de Abril 04 no vale la pena observarlo, porque si llegáramos a ver algo
sería el primer contacto, pero la Luna estará ya ocultándose.
Suerte y recuerden enviarme sus observaciones a: [email protected]
¿Cuánto de la Luna podemos ver desde la Tierra?
Jesús H. Otero A.
En un momento determinado, nunca
podemos ver más de un 50% de la superficie
lunar, pero debido al movimiento de Libración, a
lo largo del tiempo podemos observar hasta un
59% de la superficie de nuestro compañero
planetario.
Este movimiento que hace cabecear a
nuestro satélite hacia el Norte y el Sur, y hacia los
lados Este y Oeste, nos permite ver un 9% más de
la superficie lunar.
Desde casi su formación, cuando la Luna
estaba unas 10 veces más cerca de la Tierra que
ahora, el movimiento de rotación de la Luna fue
frenado por la fuerza de las mareas gravitatorias,
esto causo que se estableciera una rotación
resonante 1:1, también llamada sincrónica, esto
es, el astro de la noche tarda lo mismo en efectuar
un giro sobre sí mismo, que en girar una vez alrededor de nuestro planeta, enseñándonos por ello siempre la misma cara.
Habiendo así una cara visible y una cara oculta. Algunos dicen la cara oscura de la Luna, pero esto es un error, no hay una
cara oscura, ambos lados reciben por igual luz solar.
Si observamos la Luna
por un tiempo veremos como parecen
moverse los relieves lunares, notándose
estos un poco más al Norte o Sur, o Este
u Oeste. Esto puede notarse fácilmente
en la foto arriba, donde el cráter Tycho
pareciera estar más al Norte en la imagen
de la derecha.
Pero la Luna no solo cambia un
poquito al Este – Oeste, y Norte – Sur,
pues hay varios tipos de libraciones que
hacen a este movimiento más complejo.
Por si esto fuera poco la luna no
exhibe siempre el mismo tamaño. Como
la órbita lunar es elíptica, el tamaño
aparente de la Luna también varía si se
encuentra en Perihelio o Afelio, es decir
en el momento más cercano o más lejano de su órbita.
Solar Variability and Terrestrial Climate
Dr. Tony Phillips, NASA
Jan. 8, 2013: In the galactic scheme of things, the Sun is a remarkably constant star. While some stars exhibit
dramatic pulsations, wildly yo-yoing in size and brightness, and sometimes even exploding, the luminosity of our own
sun varies a measly 0.1% over the course of the 11-year solar cycle.
There is, however, a dawning realization among researchers that even these apparently tiny variations can have a
significant effect on terrestrial climate. A new report issued by the National Research Council (NRC), "The Effects of
Solar Variability on Earth's Climate," lays out some of the surprisingly complex ways that solar activity can make itself
felt on our planet.
These six extreme UV images of the sun, taken by NASA's Solar Dynamics Observatory, track the rising level of solar
activity as the sun ascends toward the peak of the latest 11-year sunspot cycle.
Understanding the sun-climate connection requires a breadth of expertise in fields such as plasma physics, solar
activity, atmospheric chemistry and fluid dynamics, energetic particle physics, and even terrestrial history. No single
researcher has the full range of knowledge required to solve the problem. To make progress, the NRC had to assemble
dozens of experts from many fields at a single workshop. The report summarizes their combined efforts to frame the
problem in a truly multi-disciplinary context.
One
of
the
participants, Greg Kopp of the
Laboratory for Atmospheric
and Space Physics at the
University
of
Colorado,
pointed out that while the
variations in luminosity over
the 11-year solar cycle amount
to only a tenth of a percent of
the sun's total output, such a
small
fraction
is
still
important. "Even typical short
term variations of 0.1% in
incident irradiance exceed all
other energy sources (such as
natural radioactivity in Earth's
core) combined," he says.
Of
particular
importance is the sun's
extreme ultraviolet (EUV)
radiation, which peaks during
the years around solar
maximum.
Within
the
relatively narrow band of EUV
wavelengths, the sun’s output varies not by a minuscule 0.1%, but by whopping factors of 10 or more. This can strongly
affect the chemistry and thermal structure of the upper atmosphere.
Space-borne measurements of the total solar irradiance (TSI) show ~0.1 percent variations with solar activity on
11-year and shorter timescales. These data have been corrected for calibration offsets between the various instruments
used to measure TSI. SOURCE: Courtesy of Greg Kopp, University of Colorado.
Several researchers discussed how changes in the upper atmosphere can trickle down to Earth Surface. There are
many “Top Down” pathways for the Sun´s influence. For instance Charles Jackman of the Goddard Space Flight Center
described how Nitrogen Oxides (NOx) created by solar energetic particles and cosmic rays in the stratosphere could
reduce Ozone labels by a few percent. Because Ozone absorbs UV radiation, less Ozone means that more UV rays from
the Sun would reach Earth surface.
Several researchers discussed how changes in the upper atmosphere can trickle down to Earth Surface. Isaac
Held of NOAA took this one step further. He describes how loss of Ozone in the stratosphere could alter the dynamics of
the atmosphere below it. “the cooling of polar stratosphere associated with loss of Ozone increases the horizontal
temperature gradient near the Tropopause”, he explains. “This alter the flux of angular momentum by mid-latitudes
eddies. [Angular momentum is important because] the Angular momentum budget of troposphere controls the surface
westerlies”. In other
words, solar activity
feld in the upper
atmosphere
can,
through a complicate
series of influences,
push surface storm
tracks off course.
How
incoming
galactic
cosmic rays and
solar
protons
penetrate
the
atmosphere.
SOURCE:
C.
Jackman,
NASA
Goddard
Space
Flight Center, “The Impact of Energetic Particle Precipitation on the Atmosphere,” presentation to the
Workshop on the Effects of Solar Variability on Earth’s Climate, September 9, 2011.
Many of the mechanisms proposed at the workshop had a Rube Goldberg-like quality. They relied on multi-step
interactions between multiple layers of atmosphere and ocean, some relying on chemistry to get their work done, others
leaning on thermodynamics or fluid physics. But just because something is complicated doesn't mean it's not real.
Indeed, Gerald Meehl of the National Center for Atmospheric Research (NCAR) presented persuasive evidence
that solar variability is leaving an imprint on climate, especially in the Pacific. According to the report, when researchers
look at sea surface temperature data during sunspot peak years, the tropical Pacific shows a pronounced La Nina-like
pattern, with a cooling of almost 1o C in the equatorial eastern Pacific. In addition, "there are signs of enhanced
precipitation in the
Pacific ITCZ (InterTropical Convergence
Zone ) and SPCZ (South
Pacific
Convergence
Zone) as well as abovenormal
sea-level
pressure in the midlatitude
North and
South
Pacific,"
correlated with peaks in
the sunspot cycle.
The solar cycle
signals are so strong in
the Pacific, that Meehl
and colleagues have
begun to wonder if
something in the Pacific
climate system is acting
to amplify them. "One
of
the
mysteries
regarding Earth's climate
system ... is how the
relatively
small
fluctuations of the 11year solar cycle can
produce the magnitude
of the observed climate
signals in the tropical
Pacific."
Using
supercomputer models of climate, they show that not only "top-down" but also "bottom-up" mechanisms involving
atmosphere-ocean interactions are required to amplify solar forcing at the surface of the Pacific.
Composite averages for December-January-February for peak solar years. SOURCE: G.A. Meehl,
J.M. Arblaster, K. Matthes, F. Sassi, and H. van Loon,
Amplifying the Pacific climate system response to a small 11 year solar cycle forcing, Science 325:1114-1118,
2009; reprinted with permission from AAAS.
In recent years, researchers have considered the possibility that the sun plays a role in global warming. After all,
the sun is the main source of heat for our planet. The NRC report suggests, however, that the influence of solar variability
is more regional than global. The Pacific region is only one example.
Caspar Amman of NCAR noted in the report that "When Earth's radiative balance is altered, as in the case of a
change in solar cycle forcing, not all locations are affected equally. The equatorial central Pacific is generally cooler, the
runoff from rivers in Peru is reduced, and drier conditions affect the western USA."
Raymond Bradley of UMass, who has studied historical records of solar activity imprinted by radioisotopes in
tree rings and ice cores, says that regional rainfall seems to be more affected than temperature. "If there is indeed a solar
effect on climate, it is manifested by changes in general circulation rather than in a direct temperature signal." This fits in
with the conclusion of the IPCC and previous NRC reports that solar variability is NOT the cause of global warming over
the last 50 years.
Much has been made of the probable connection between the Maunder Minimum, a 70-year deficit of sunspots
in the late 17th-early 18th century, and the coldest part of the Little Ice Age, during which Europe and North America
were subjected to bitterly cold winters. The mechanism for that regional cooling could have been a drop in the sun’s
EUV output; this is, however, speculative.
The yearly averaged sunspot number for a period of 400 years (1610-2010). SOURCE: Courtesy of NASA
Marshall Space Flight Center.
Dan Lubin of the Scripps Institution of Oceanography pointed out the value of looking at sun-like stars
elsewhere in the Milky Way to determine the frequency of similar grand minima. “Early estimates of grand minimum
frequency in solar-type stars ranged from 10% to 30%, implying the sun’s influence could be overpowering. More recent
studies using data from Hipparcos (a European Space Agency astrometry satellite) and properly accounting for the
metallicity of the stars, place the estimate in the range of less than 3%.” This is not a large number, but it is significant.
Indeed, the sun could be on the threshold of a mini-Maunder event right now. Ongoing Solar Cycle 24 is the
weakest in more than 50 years. Moreover, there is (controversial) evidence of a long-term weakening trend in the
magnetic field strength of sunspots. Matt Penn and William Livingston of the National Solar Observatory predict that by
the time Solar Cycle 25 arrives, magnetic fields on the sun will be so weak that few if any sunspots will be formed.
Independent lines of research involving helio seismology and surface polar fields tend to support their conclusion. (Note:
Penn and Livingston were not participants at the NRC workshop.)
“If the sun really is entering an unfamiliar phase of the solar cycle, then we must redouble our efforts to understand the
sun-climate link,” notes Lika Guhathakurta of NASA’s living with a Star Program, which helped fund the NRC study.
“The report offers some good ideas for how to get started.”
This image of the Sun's upper photosphere shows bright and dark magnetic structures responsible for variations
in TSI. SOURCE: Courtesy of P. Foukal, Heliophysics, Inc.
In a concluding panel discussion, the researchers identified a number of possible next steps. Foremost among
them was the deployment of a radiometric imager. Devices currently used to measure total solar irradiance (TSI) reduce
the entire sun to a single number: the total luminosity summed over all latitudes, longitudes, and wavelengths. This
integrated value becomes a solitary point in a time series tracking the sun’s output.
In fact, as Peter Foukal of Heliophysics, Inc., pointed out, the situation is more complex. The sun is not a
featureless ball of uniform luminosity. Instead, the solar disk is dotted by the dark cores of sunspots and splashed with
bright magnetic froth known as faculae. Radiometric imaging would, essentially, map the surface of the sun and reveal
the contributions of each to the sun’s luminosity. Of
particular interest are the faculae. While dark sunspots tend
to vanish during solar minima, the bright faculae do not.
This may be why paleo climate records of sun-sensitive
isotopes C-14 and Be-10 show a faint 11-year cycle at work
even during the Maunder Minimum. A radiometric imager,
deployed on some future space observatory, would allow
researchers to develop the understanding they need to project
the sun-climate link into a future of prolonged spotlessness.
Some attendees stressed the need to put sun-climate
data in standard formats and make them widely available for
multidisciplinary study. Because the mechanisms for the
sun’s influence on climate are complicated, researchers from
many fields will have to work together to successfully model
them and compare competing results. Continued and
improved collaboration between NASA, NOAA and the NSF
are keys to this process.
Hal Maring, a climate scientist at NASA
headquarters who has studied the report, notes that “lots of
interesting possibilities were suggested by the panelists.
However, few, if any, have been quantified to the point that we can definitively assess their impact on climate.”
Hardening the possibilities into concrete, physically-complete models is a key challenge for the researchers.
Finally, many participants noted the difficulty in deciphering the sun-climate link from paleo climate records
such as tree rings and ice cores. Variations in Earth’s magnetic field and atmospheric circulation can affect the deposition
of radioisotopes far more than actual solar activity. A better long-term record of the sun’s irradiance might be encoded in
the rocks and sediments of the Moon or Mars. Studying other worlds
might hold the key to our
own.
The full report,
“The Effects of Solar
Variability on Earth’s
Climate,” is available
from
the
National
Academies Press at
http://www.nap.edu/catal
og.php?record_id=13519
.
¿Pudo el Bosón de Higgs hacer colapsar el Universo?
Cosmólogos británicos han concluido que el Universo pudo no haber durado más de un segundo
luego del Big Bang.
Foto: El Telescopio BICEP 2 en
uno de los dos atardeceres que
ocurren en el año en el Polo Sur.
El observatorio MAPO (hogar de
la Red de telescopios Keck), y la
estación del Polo Sur se pueden
observar en el fondo.
Los cosmólogos británicos están
confundidos, ellos predijeron que
nuestro Universo no debió durar
más de un segundo. Esta extraña
conclusión es el resultado de
combinar
las
últimas
observaciones del cielo con el
reciente descubrimiento del
Bosón de Higgs. El Dr. Robert
Hogan, del King’s College
London (KCL), presentó el
trabajo en Junio 24, 2014. En la
reunión de la Royal Sociedad
Nacional
de
Astronomía
Astronomy, en Portsmouth.
Después que nuestro Universo empezó como el Big Bang, se cree que tuvo un corto período de rápida expansión
que conocemos como Inflación Cósmica. Aunque algunos detalles de este proceso no son bien entendidos, los
cosmólogos han sido capaces de hacer predicciones sobre cómo este proceso afecta al Universo que vemos hoy día.
En Marzo de 2014, investigadores colaboradores del BICEP 2 dijeron que habían detectado uno de los efectos
predichos. Si es verdad, estos resultados son un gran avance en nuestra comprensión de la Cosmología y confirmación de
la Teoría de la Inflación, pero esto ha sido muy controversial y no totalmente aceptado por los cosmólogos.
En el estudio, científicos del KCL han investigado lo que las observaciones del BICEP 2 significan para la
estabilidad del Universo. Para hacerlo combinaron los resultados con avances recientes en la Física de Partículas. La
detección del Bosón de Higgs en el Gran colisionador de Hadrones, anunciado en Julio de 2012: desde entonces se ha
aprendido mucho sobre sus propiedades.
Medidas del Bosón de Higgs han permitido a los Físicos de Partículas que nuestro Universo se asienta en un
“Valle del Campo de Higgs”, el cual describe la manera que otras partículas poseen masa. Sin embargo hay un “Valle”
diferente que es mucho más profundo, pero nuestro Universo no cae allí debido a una gran barrera energética.
El problema es que los resultados del BICEP 2 predicen que nuestro Universo ha recibido potentes impulsos
durante la fase de Inflación, empujándolo al otro “Valle” del Campo de Higgs en una fracción de segundo. Si esto hubiera
ocurrido, el Universo habría colapsado en un instante.
Robert Hogan, líder del estudio dice que esto es una predicción inaceptable, pues si esto hubiera ocurrido, no
estaríamos aquí discutiéndolo. Tal vez los resultados del BICEP 2 contienen un error. Si no, debe haber otros procesos,
aun desconocidos, que previnieron que el Universo colapsara. Si los resultados del BICEP 2 son correctos, entonces esto
nos dice que existe una nueva e interesante Física de Partículas más allá del modelo estándar.
NOTA: Ver:
Mensajero 71, Primera Evidencia de la Inflación
Mensajero 68, Física de Partículas
Mensajero 67, Hawkins y el Origen del Universo
Mensajero 66, Bosón de Higgs u otra partícula
Mensajero 64, El Bosón de Higgs
Ocultación de Marte por la Luna, Julio 06, 2014.
Observador: Jorge Luis Salas m / ACA (Asociación Carabobeña de
Astronomía) sitio web: 114milimetros.blogspot.com
Longitud: -67.960280555556
Latitud: 10.261688888889
Altura: +498.00 metros
Huso horario: UTC-4.5
Contacto
Tiempo Estimado
Tiempo Estimado
(Velocidad de la luz infinita) (Velocidad de la luz finita)
Contacto 1 (Marte toca la Luna)
2014 JUL 05 22:10:55.84
2014 JUL 05 22:10:45.70
Contacto 1 punto medio (La mitad de Marte esta ocultada) 2014 JUL 05 22:11:51.34
2014 JUL 05 22:11:43.12
Contacto 2 (Marte esta ocultado completamente)
2014 JUL 05 22:12:53.26
2014 JUL 05 22:12:46.99
Contacto 3 (Marte empieza a emerger)
2014 JUL 05 22:27:52.57
2014 JUL 05 22:27:14.77
Contacto 3 punto medio (La mitad de Marte ha emergido) 2014 JUL 05 22:28:51.51
2014 JUL 05 22:28:16.29
Contacto 4 (Fin de la ocultación)
2014 JUL 05 22:29:10.78
2014 JUL 05 22:29:45.15
Duración del evento:+0:16:33.164 horas
Trayectoria de Marte Respecto a la Luna:
Tiempos de contacto para la latitud local, nótese
que será una ocultación rasante. Desapareciendo en T1 Y
reapareciendo en T2
Ubicación: San Diego, Carabobo, Venezuela
Latitud: N 10° 15’ 42.08’’
Longitud: O 67°57’37.01’’
Altitud (msnm) : 498
Equipo utilizado:
Binocular TASCO 10X70
Telescopio reflector newtoniano Celestron Firstscope 76/300
Oculares: 4mm 9mm 20mm 32 mm
Filtro lunar Celestron
Procedimiento: Visual con binocular y Telescopio
Software/ application:
GPS STATUS, Time the sat,
cronómetro
Datos de Observación :
Contacto
TC. HLV (T.U. – 4,5h)
C-1 Marte toca la Luna
22:10:11,778
C 2 ½ Marte ocultado
22:11:27,340
C-3 Marte Ocultado
22:12:39,000
C-4 Marte Sale
22:27:34,795
C-5 ½ Marte emerge
22:28:59,000
C-6 Fin de la Ocultación
22:30:14,540
En Caracas estuvo nublado y no se realizaron
observaciones, solo Carlos Quintana pudo medir
La gran Mancha Roja de Júpiter está más pequeña que nunca antes
NASA
La Mancha Roja, un ícono del planeta Júpiter, la cual es un sistema de alta presión circular más grande que
nuestro planeta, se ha empequeñecido al menor tamaño jamás medido.
Imagen: NASA, ESA, and A. Simon (Goddard Space Flight Center)
Los astrónomos han seguido el empequeñecimiento de la Gran Mancha Roja de Júpiter desde la década de los
años 30. Mediciones recientes realizadas por el Telescopio Espacial Hubble confirman que esta tiene ahora unas 10.250
millas de diámetro, el tamaño más pequeño jamás medido, según Amy Simon del Goddard Space Flight Center de la
NASA en Greenbelt, Md.
Mediciones históricas que van para atrás hasta los últimos años del siglo XIX (1.800´s), muestran una Gran
Mancha Roja de hasta 25.500 millas de diámetro en su eje mayor. Las sondas Voyager 1 y Voyager 2 midieron un
diámetro de 14.500 millas en 1979.
Comenzando en el 2012, observaciones de aficionados revelaron un notable incremento en la taza de
encogimiento. El talle de la GMR se está encogiendo 580 millas por año, y el óvalo de la GMR paso a ser un circulo. La
causa del encogimiento aún no se ha explicado.
En observaciones recientes se han observado pequeños remolinos se están alimentando de la tormenta. Se piensa
que ellas son responsables de los cambios acelerados por alteración de la dinámica y energía interna de la Gran Mancha
Roja.
El equipo de astrónomos liderizado por A. Simon, planea estudiar los pequeños remolinos y estudiar la dinámica
de la GMR para determinar si ellos pueden alimentar o frenar el momentum al entrar en el vórtice.
Comparaciones realizadas con el Hubble tomadas en 1995, cuando el eje mayor de la GMR era 13.020 millas,
con mediciones del 2009, muestran que en este año medía 11.130 millas. Unas 1890 millas menos.
Pero Júpiter es un planeta muy activo en su atmósfera, hace 3 años la Banda Ecuatorial Sur desapareció u estuvo
2 años ausente. Algo así podría estar pasando a la Gran Mancha Roja, o tal vez este rasgo distintivo del planeta este
llegando a su fin. Aún hay mucho que investigar.
La acidificación actual del mar es mucho más rápida que la de hace 56 millones de años
La acidificación actual del mar es mucho más rápida que la de hace 56 millones de años
Hace unos 56 millones de años, hubo un período de calentamiento global abrupto, el cual se conoce como el Máximo
Térmico del Paleoceno-Eoceno (MTPE, o PETM por sus siglas en inglés). Durante esta etapa geológica, un pulso masivo
de dióxido de carbono emitido hacia la atmósfera elevó ostensiblemente las temperaturas a escala global. En los océanos,
los sedimentos del carbonato se disolvieron, algunas especies se extinguieron y otras experimentaron un fuerte cambio de
rumbo evolutivo.
Lejos de ser un fenómeno de interés exclusivo para los estudiosos del pasado, el Máximo Térmico del
Paleoceno-Eoceno es hoy en día un tema de la máxima actualidad, ya que cada vez está más claro que se trata del análogo
más cercano, por similitud y por cercanía en el tiempo, al actual calentamiento global.
Entre los efectos comunes a ambos episodios figura la acidificación oceánica. La comunidad científica ha sospechado
desde hace mucho tiempo que fue la acidificación oceánica la ejecutora de los cambios nocivos en el mar que
perjudicaron a los antiguos ecosistemas marinos. Aquella crisis medioambiental aparece marcada claramente en los
registros fósiles y geológicos.
De manera similar a lo que ocurre hoy, la creciente abundancia del dióxido de carbono propició que éste se
combinase con el agua salada de los océanos de tal modo que alteró las propiedades químicas de ésta.
Ahora unos científicos han logrado cuantificar por primera vez la magnitud de la acidificación de la superficie oceánica
durante el Máximo Térmico del Paleoceno-Eoceno, y las noticias no son buenas: Nuestros océanos actuales están en
camino de acidificarse tanto o más que en aquel entonces, sólo que a una velocidad mucho más rápida, que puede ser
hasta 10 veces más veloz que en esa época de referencia.
[Img #20956]
Los foraminíferos de la especie Aragonia velascoensis se extinguieron, junto con otras criaturas marinas, hace
unos 56 millones de años, por culpa de la acidificación oceánica, rápida para lo que el ritmo de la evolución es capaz de
afrontar, pero que pese a todo fue mucho más lenta que la actual. (Foto: Ellen Thomas / Universidad Yale)
El equipo de la paleoceanógrafa Bärbel Hönisch, del Observatorio Terrestre Lamont-Doherty, adscrito a la Universidad
de Columbia, en la ciudad de Nueva York, y Ellen Thomas, de la Universidad Yale en New Haven, Connecticut, todas
estas entidades en Estados Unidos, estima que la acidez oceánica aumentó en aproximadamente un 100 por cien a lo largo
de un periodo de unos mil años o más, y se quedó así durante los siguientes 70.000 años. En este ambiente alterado
radicalmente, algunas especies se extinguieron inexorablemente mientras que otras se adaptaron y evolucionaron.
Los océanos, cual héroes silenciosos de nuestros tiempos, han absorbido cerca de un tercio de las emisiones de carbono
que los humanos hemos bombeado a la atmósfera desde la industrialización. Con su acción protectora, han ayudado a
mantener la temperatura más baja de lo que habría ya llegado a ser en su ausencia. Pero esa captura del carbono tiene su
precio. Las reacciones químicas causadas por ese exceso de CO2 han hecho que el agua de mar sea más ácida,
desposeyéndola de los iones de carbonato que corales, moluscos y algunas especies de plancton necesitan para desarrollar
sus conchas y esqueletos.
En los últimos 150 años, el pH de los océanos ha descendido de manera significativa (o sea que su agua se ha
vuelto más ácida). Se estima que desde
ahora y hasta finales de este siglo, la
caída del pH oceánico será incluso
mayor que la registrada en el último siglo
y medio. Sumando la caída de los
últimos 150 años con la pronosticada
para el siglo actual, el aumento de acidez
marina resultante es un poco mayor que
el estimado para todo el Máximo
Térmico del Paleoceno-Eoceno. Lo más
inquietante, sin embargo, es que,
mientras que el cambio de pH en el
Máximo Térmico del Paleoceno-Eoceno
se obró a lo largo de unos mil años, el
actual cambio de pH, si se cumplen las
previsiones, se habrá obrado en un
periodo mucho menor, de tan solo unos
250 años.
El Spitzer de NASA, WISE Encuentra un Sol vecino cercano y frío.
NASA
Esta concepción artística muestra
al
objeto
llamado
WISE
J085510.83-071442.5, la más fría
enana marrón conocida. Estas
estrellas son pequeños cuerpos
parecidos a estrellas que carecen
de masa para quemar sus
combustibles nucleares.
El NASA's Wide-field
Infrared Survey Explorer (WISE) y
el Spitzer Space Telescope han
descubierto lo que parece ser la más
fría estrella enana marrón conocida,
una estrella muy tenue que
sorprendentemente es tan fría como
los polos terrestres.
Imágenes de telescopios espaciales también descubrieron que su distancia es de solo 7.2 años luz, lo que la
coloca entre los 4 sistemas estelares más cercanos a nuestro Sol. El sistema más cercano es un trío de estrellas que
llamamos Alfa Centauro y que dista a solo 4.3 años luz de nosotros.
"Es muy excitante descubrir un nuevo vecino tan cercano a nuestro Sistema Solar”, dice Kevin Luhman, un
astrónomo de la Universidad de la Pennsyvania State Universitys University Park Center para Exoplanetas y Mundos
Habitable, "Y dada su temperatura extrema, nos puede decir mucho sobre las atmósferas planetarias, las cuales poseen
frecuentemente, temperaturas similares”.
Las estrellas Enanas Marrones comienzan sus vidas como estrellas, bolas de gas que colapsan, pero por no
poseer masa suficiente para encender sus hornos nucleares no pueden radiar energía y brillar como estrellas. La nueva
estrella Enana Marrón, la más fría jamás descubierta es llamada: WISE J085510.83-071442.5. Ella tiene una fría
temperatura entre -54 y 9º Fahrenheit, (-48 y 9º Celsius). Los records anteriores para la enana marrón más fría, también
descubierta por el WISE y el Spitzer, era como la temperatura de una habitación normal.
El WISE fue capaz de captar el raro objeto por realizar dos surveys del cielo en infrarrojo, observando áreas
hasta 3 veces. Objetos fríos como las enanas marrones pueden ser invisibles al observarlas con telescopios de luz visible,
pero su brillo térmico aparece en IR aunque sean objetos fríos. En adición, mientras más cercano sea un objeto, más
parecerá moverse en imágenes tomadas tras varios meses.
Este objeto se movía mucho en los datos del WISE, lo que nos dijo que era muy cercano.
Luego de notar el rápido movimiento del WISE J085510.83-071442.5 en Marzo de 2013, Luhman analizó
imágenes de Spitzer y el Telescopio Géminis Sur en cerro Pachón, chile. Las observaciones del telescopio infrarrojo
Spitzer, ayudaron a determinar la helada temperatura de la enana roja. Combinando las observaciones de WISE y Spitzer,
tomadas en posiciones diferentes alrededor del Sol, realizaron la medición de distancias por efecto de la Paralaje. Este es
el mismo principio que explica el movimiento de un dedo de su mano, cuando lo coloca frente a su rostro y lo mira con el
ojo derecho y luego el izquierdo.
Es interesante que después de décadas estudiando el cielo, aún no poseamos in inventario completo de los
vecinos más cercanos al Sol, de acuerdo a lo dicho por Michael Werner, científico del proyecto Spitzer, del Jet Propulsion
Laboratory, de NASA, en Pasadena, California. El JPL gerencia y maneja el Spitzer. Este nuevo resultado es excitante,
pues demuestra el poder de la exploración astronómica utilizando nuevas herramientas como los telescopios IR Wise y
Spitzer.
El WISE J085510.83-071442.5 se estima que tiene entre 3 y 10 masas de Júpiter. Con esta masa debe ser un
gigante gaseoso igual al planeta, que fue eyectado de su sistema estelar. Algunos científicos creen sin embargo que se
trata de un estrella Enana Marrón y no de un planeta, pues se sabe que estas son muy comunes. Si es así es una de las
Enanas Marrones con menos masa conocida.
En Marzo del 2013, los análisis imágenes de Luhman realizados con el WISE descubrió un par de enanas
marrones más calientes, a una distancia de 6.5 Años Luz, haciendo a este sistema el tercero más cercano al Sol. Su
búsqueda de objetos rápidos también demostró que el sistema solar exterior probablemente no posee un planeta grande no
descubierto y al cual se ha llamado Planeta X o Némesis.
A New View of the Red Planet
Authors: Tenielle Gaither, [email protected]; Kenneth Tanaka, [email protected]; James Skinner,
[email protected]; Jennifer LaVista, [email protected]
Get ready, because now you can explore the most comprehensive representation of Mars with a new global
geologic map created by the U.S. Geological Survey. This new view of the “Red Planet’s” surface provides a framework
for continued scientific investigation of Mars as the long-range target for human space exploration.
What Does the New Map Show?
The USGS-led mapping effort reveals that the Martian surface is generally older than previously thought. Three
times as much surface area dates to the first major geologic time period – the Early Noachian Epoch – than was
previously mapped. This timeframe is the earliest part of the Noachian Period, which ranges from about 4.1 to about 3.7
billion years ago, and was characterized by high rates of meteorite impacts, widespread erosion of the Martian surface
and the likely presence of abundant surface water.
The map also confirms previous work that suggests Mars had been geologically active until the present day.
There is evidence that major changes in Mars’ global climate supported the temporary presence of surface water and
near-surface groundwater and ice. These changes were likely responsible for many of the major shifts in the environments
where Martian rocks were formed and subsequently eroded. This new map will serve as a key reference for the origin,
age and historic change of geological materials anywhere on Mars.
Why Explore Space?
For hundreds of years, geologic maps have helped drive scientific thought. This new global geologic map of
Mars, as well as the recent global geologic maps of Jupiter’s moons Ganymede and Io, also illustrates the overall
importance of geologic mapping as an essential tool for the exploration of the solar system.
“Spacecraft exploration of Mars over the past couple decades has greatly improved our understanding of what
geologic materials, events and processes shaped its surface,” said USGS scientist and lead author, Dr. Kenneth Tanaka.
“The new geologic map brings this research together into a holistic context that helps to illuminate key relationships in
space and time, providing information to generate and test new hypotheses.”
Out of this World Science Takes Time
The new map brings together observations and scientific findings from four orbiting spacecraft that have been
acquiring data for more than 16 years. The result is an updated understanding of the geologic history of the surface of
Mars – the solar system’s most Earth-like planet and the only other one in our Sun’s “habitable zone.”
The Martian surface has been the subject of scientific observation since the 1600s, first by Earth-based
telescopes, and later by fly-by missions and orbiting spacecraft. The Mariner 9 and Viking Orbiter missions produced the
first planet-wide views of Mars’ surface, enabling publication of the first global geologic maps (in 1978 and 1986-87,
respectively) of a planetary surface other than the Earth and the Moon. A new generation of sophisticated scientific
instruments flown on the Mars Global Surveyor, Mars Odyssey, Mars Express and Mars Reconnaissance Orbiter
spacecraft has provided diverse, high quality data sets that enable more sophisticated remapping of the global-scale
geology of Mars.
How the USGS Got Involved in Space Science
The production of planetary cartographic products has been a focal point of research at the USGS Astrogeology
Science Center since its inception in the early 1960s. The USGS began producing planetary maps in support of the Apollo
Moon landings, and continues to help establish a framework for integrating and comparing past and future studies of
extraterrestrial surfaces. In many cases, these planetary geologic maps show that, despite the many differences between
bodies in our solar system, there are many notable similarities that link the evolution and fate of our planetary system
together.
The mission of the USGS Astrogeology Science Center is to serve the nation, the international planetary science
community and the general public’s pursuit of new knowledge of our solar system. The team’s vision is to be a national
resource for the integration of planetary geosciences, cartography and remote sensing. As explorers and surveyors with a
unique heritage of proven expertise and international leadership, USGS astrogeologists enable the ongoing successful
investigation of the solar system for humankind.
Enabling Future Exploration
“Findings from the map will enable researchers to evaluate potential landing sites for future Mars missions that may
contribute to further understanding of the planet’s history,” said USGS Acting Director Suzette Kimball. “The new Mars
global geologic map will provide geologic context for regional and local scientific investigations for many years to
come.”
The project was funded by NASA through its Planetary Geology and Geophysics Program.
Violenta historia del Sol joven resuelve misterio de los Meteoritos
ESA
Un grupo de astrónomos ha empleado el telescopio espacial
Herschel de ESA para estudiar los violentos comienzos de una
estrella tipo Sol, encontrando indicios de potentes vientos estelares
que podrían resolver un extraño misterio sobre meteoritos.
A pesar de su tranquila apariencia en el cielo nocturno, las
estrellas son hornos abrasadores que llegan a la vida a través de
procesos tumultuosos, y nuestro Sol, de 4500 millones de años de
edad, no es una excepción. Para conocer un poco más sobre sus
duros inicios, los astrónomos recogen pistas, no sólo en el Sistema
Solar, sino también estudiando estrellas jóvenes en otros lugares de
nuestra Galaxia.
Empleando Herschel para estudiar la composición química
de regiones donde las estrellas están naciendo hoy en día, un equipo
de astrónomos ha observado que un objeto en particular es
diferente. La fuente inusual es un prolífico vivero estelar llamado
OMC2 FIR4, una agrupación de estrellas nuevas situadas en el
interior de una nube gaseosa y polvorienta, cerca de la famosa
Nebulosa de Orión.
"Para nuestra sorpresa, descubrimos que la proporción
entre dos especies químicas, una basada en el carbono y oxígeno y
la otra en el nitrógeno, es mucho más pequeña en este objeto que en
cualquier otra protoestrella que conozcamos", afirma la Dra. Cecilia
Ceccarelli, quien dirigió el estudio."La causa más probable en este
ambiente es un violento viento de partículas muy energéticas,
expulsado por lo menos por una de las estrellas embrionarias que
están tomando forma en este huevo protoestelar", añade la Dra.
Ceccarelli.
Los astrónomos piensan que un viento violento parecido
de partículas también barrió el Sistema Solar primitivo, y este
descubrimiento podría finalmente constituir una explicación al origen de un elemento químico particular observado en
meteoritos, el berilio 10.
Junio fue el mes más cálido en la Tierra desde 1880
NOAA
Imagen de la Tierra desde la Estación Espacial
Internacional
La, Administración Nacional Atmosférica
y Oceánica de EE UU registró que la media del
planeta se colocó en 15,5°C.
El mes de junio registró las temperaturas
globales más cálidas desde que comenzaron los
registros en 1880, al superar la media de 15,5°C
(59,9°F) por 0,72°C (1,30°F), informó hoy la
Administración Nacional Atmosférica y Oceánica
de EEUU, (NOAA, por sus siglas en Inglés), en su
reporte mensual. Tanto la temperatura de la
superficie terrestre como la de los océanos
alcanzaron temperaturas superiores a la media.
No obstante, Jessica Blunden, científica de la NOAA, apuntó que "el calentamiento fue impulsado por las
temperaturas récord en el océano", y agregó que parte de esta subida se debió al comienzo del fenómeno de El Niño, el
calentamiento de las aguas tropicales del Pacífico.
La de la tierra fue 0,95°C (1,71°F) mayor que la media de 13,3°C (55,9°F), y se situó como el séptimo junio más
cálido; mientras que la del océano fue de 0,64°C (1,15°F) por encima de la media de 16,4°C (61,5°F), y se convierte así
en el junio más cálido desde que se empezaron a compilar datos.
Estos datos reflejan también un repunte en las temperaturas globales en los primeros seis meses del año: las
combinadas de tierra y mar fueron 0,67°C (1,21°F) superiores a la media de 13,5°C (56,3°F), las terceras más altas para
un período enero-junio desde 1880.
Esta alza en las temperaturas se produjo de manera general en todo el mundo, ya que se batieron récords en
Groenlandia, el norte de Sudamérica, el este y centro de África y el sudeste asiático, así como en Nueva Zelanda.
La agencia federal de EE.UU. recordó en su reporte mensual que nueve de los diez meses de junio más cálidos
registrados han tenido lugar en el siglo XXI.
Asimismo, indicó que el último mes de junio por debajo de la media se produjo en 1976.
Sonda Voyager 1 podría no haber alcanzado el espacio interestelar.
En 2012, el equipo de la misión Voyager anunció que la nave Voyager 1 había pasado al espacio interestelar,
viajando más lejos de lo que lo ha hecho cualquier objeto de fabricación humana. Pero en los casi dos años que han
transcurrido desde ese anuncio histórico, y a pesar de las observaciones posteriores que lo respaldaban, continúa la
incertidumbre acerca de si la Voyager 1 realmente cruzó la frontera. Hay algunos científicos que dicen que la nave
espacial todavía se encuentra dentro de la heliosfera (la región del espacio dominada por el Sol y su viento de partículas
energéticas) y que aún no ha alcanzado el espacio entre las estrellas.
Ahora, dos científicos del equipo de la Voyager han
desarrollado una prueba que podría demostrar de una vez por
todas si la Voyager 1 ha cruzado la frontera. Los científicos
predicen que en los próximos dos años Voyager 1 cruzará la
capa de corriente eléctrica (la superficie dentro de la heliosfera
donde la polaridad del campo magnético del Sol cambia de
positiva a negativa). La nave detectará una inversión del campo
magnético, demostrando que todavía se encuentra dentro de la
heliosfera. Pero si la inversión del campo magnético no se
produce dentro de un año o dos, tal como se espera, eso
confirmaría que Voyager 1 ya ha pasado al espacio interestelar.
Las naves espaciales Voyager 1 y 2 fueron lanzadas en
1977 hacia Júpiter y Saturno. Desde entonces la misión se ha
extendido a la exploración de los límites más exteriores de la
influencia del Sol y aún más allá. Voyager 2, que también pasó
por Urano y Neptuno, está de camino al espacio interestelar.
La supernova más brillante de la historia
A 7.000 años luz de la Tierra, era tan espectacular que pudo ser contemplada durante más de tres años en el siglo
XI. Ahora, los científicos saben qué la provocó
NASA/NRAO/Middlebury College
El remanente de supernova, a 7.000 años
luz de la Tierra
Una investigación, en la que ha
participado el Consejo Superior de
Investigaciones Científicas (CSIC), ha
descubierto el origen del que hasta ahora
se considera el "evento estelar más
brillante" que ha podido ser contemplado
en la historia desde la Tierra, la supernova
SN1006, que tuvo lugar en el año 1006 a
unos 7.000 años luz de la Tierra, fruto de
la fusión de dos enanas blancas, según ha
publicado la revista Nature en su portada.
De esta forma, el CSIC señala que
este evento estelar se clasifica dentro de las
supernovas de tipo Ia, que son aquellas
generadas por sistemas binarios en los que
dos objetos astronómicos están ligados
entre sí por su fuerza gravitatoria.
Asimismo, apunta que el estudio
calcula que la luz emitida por SN1006 fue
equivalente a "una cuarta parte" de la del
brillo de la Luna, lo que respaldaría los
registros históricos de astrólogos de la
época que indican que la explosión fue visible en distintas partes del mundo durante "más de tres años" y que fue
"aproximadamente" tres veces más brillante que Venus.
Por otro lado, explica que "usualmente" estos sistemas suelen estar formados por una enana blanca y una estrella
normal que le aporta la materia necesaria para alcanzar la "masa crítica" de 1,4 veces la del Sol y, una vez alcanzada, la
enana blanca comienza la fusión de su núcleo que origina una explosión termonuclear. No obstante, ha apuntado que
"también existe la posibilidad de que la supernova se origine a causa de la fusión de dos enanas blancas conectadas entre
sí".
Por su parte, la investigadora del Instituto de Física Fundamental del CSIC Pilar Ruiz-Lapuente, que ha
participado en este estudio, ha manifestado que "la exploración en torno al lugar donde se produjo la supernova SN1006
no ha detectado a ningún candidato a compañero de la enana blanca original, lo que invita a pensar que probablemente se
produjo mediante la fusión de dos enanas blancas conectadas entre sí". Ante esto, el investigador del Instituto de
Astrofísica de Canarias Jonay González, que ha liderado el trabajo, ha argumentado que "existen tres tipos de estrellas en
la región donde tuvo lugar la explosión, las gigantes, sub gigantes y enanas, pero las observaciones sólo detectaron cuatro
estrellas gigantes situadas a la misma distancia que el remanente de la supernova".
Sin dejar pistas
Así, ha planteado que "las simulaciones numéricas no predicen a una compañera de estas características, las
cualidades de una posible estrella compañera". En este sentido, Ruiz-Lapuente ha indicado que "tras la explosión de la
supernova, la estrella compañera de la enana blanca se asemejaría más a una estrella de helio, pero ninguna de este tipo
fue detectada en la región de estudio por lo que se desprende que el origen de SN1006 tuvo lugar en la colisión de dos
enanas blancas, cuyo material fue expulsado sin dejar ningún testigo de la explosión".
Por último, la investigadora del CSIC ha apuntado que "hasta la fecha se habían encontrado algunas supernovas
extra galácticas que no mostraban ninguna señal de la existencia de la estrella compañera". Por ello, considera que estos
"nuevos resultados, junto con otros anteriores, suponen que la fusión de enanas blancas podría ser una vía usual para dar
lugar a estas violentas explosiones termonucleares". En el año 2004, Ruiz-Lapuente ya dirigió la investigación para
descubrir el origen de la supernova del año 1572, donde hallaron la estrella que acompañó a la enana blanca que provoco
este evento estelar.
Near- and Mid-Infrared Photometry of the Pleiades and a New List of Substellar Candidate Members1,2
John R. Stauffer
Spitzer Science Center, Caltech 314-6, Pasadena, CA 91125; [email protected]
Lee W. Hartmann
Astronomy Department, University of Michigan
Giovanni G. Fazio , Lori E. Allen , and Brian M. Patten
Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138
Patrick J. Lowrance , Robert L. Hurt , and Luisa M. Rebull
Spitzer Science Center, Caltech, Pasadena, CA 91125
Roc M. Cutri and Solange V. Ramirez
Infrared Processing and Analysis Center, Caltech 220-6, Pasadena, CA 91125
3
Erick T. Young , George H. Rieke , Nadya I. Gorlova , and James C. Muzerolle
Steward Observatory, University of Arizona, Tucson, AZ 85726
Cathy L. Slesnick
Astronomy Department, Caltech, Pasadena, CA 91125
Michael F. Skrutskie
Astronomy Department, University of Virginia, Charlottesville, VA 22903
ABSTRACT
We make use of new near- and mid-IR photometry of the Pleiades cluster in order to help identify proposed
cluster members. We also use the new photometry with previously published photometry to define the single-star mainsequence locus at the age of the Pleiades in a variety of color-magnitude planes. The new near- and mid-IR photometry
extend effectively 2 mag deeper than the 2MASS All-Sky Point Source catalog, and hence allow us to select a new set of
candidate very low-mass and substellar mass members of the Pleiades in the central square degree of the cluster. We
identify 42 new candidate members fainter than Ks = 14 (corresponding to 0.1 M ). These candidate members should
eventually allow a better estimate of the cluster mass function to be made down to of order 0.04 M . We also use new
IRAC data, in particular the images obtained at 8 m, in order to comment briefly on interstellar dust in and near the
Pleiades. We confirm, as expected, that—with one exception—a sample of low-mass stars recently identified as having 24
m excesses due to debris disks do not have significant excesses at IRAC wavelengths. However, evidence is also
presented that several of the Pleiades high-mass stars are found to be impacting with local condensations of the molecular
cloud that is passing through the Pleiades at the current epoch.
Subject headings: open clusters and associations: individual (Pleiades); stars: low-mass, brown dwarfs
Online material: color figure, machine-readable tables
1
This work is based (in part) on observations made with the Spitzer Space Telescope, which is operated by the Jet
Propulsion
Laboratory,
California
Institute
of
Technology,
under
NASA
contract
1407.
2
This publication makes use of data products from the Two Micron All Sky Survey, which is a joint project of the
University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded
by the National Aeronautics and Space Administration and the National Science Foundation.
3
Current address: University of Florida, 211 Bryant Space Center, Gainesville, FL 32611.
1. INTRODUCTION
Because of its proximity, youth, richness, and location in the northern hemisphere, the Pleiades has long been a
favorite target of observers. The Pleiades was one of the first open clusters to have members identified via their common
proper motion (Trumpler 1921), and the cluster has since then been the subject of more than a dozen proper-motion
studies. Some of the earliest photoelectric photometry was for members of the Pleiades (Cummings 1921), and the cluster
has been the subject of dozens of papers providing additional optical photometry of its members. The youth and nearness
of the Pleiades make it a particularly attractive target for identifying its substellar population, and it was the first open
cluster studied for those purposes (Jameson & Skillen 1989; Stauffer et al. 1989). More than 20 papers have been
subsequently published, identifying additional substellar candidate members of the Pleiades or studying their properties.
We have three primary goals for this paper. First, while extensive optical photometry for Pleiades members is
available in the literature, photometry in the near- and mid-IR is relatively spotty. We will remedy this situation by using
new 2MASS JHKs and Spitzer Infrared Array Camera (IRAC) photometry for a large number of Pleiades members. We
will use these data to help identify cluster nonmembers and to define the single-star locus in color-magnitude diagrams for
stars of 100 Myr age. Second, we will use our new IR imaging photometry of the center of the Pleiades to identify a new
set of candidate substellar members of the cluster, extending down to stars expected to have masses of order 0.04 M .
Third, we will use the IRAC data to briefly comment on the presence of circumstellar debris disks in the Pleiades and the
interaction of the Pleiades stars with the molecular cloud that is currently passing through the cluster.
In order to make best use of the IR imaging data, we will begin with a necessary digression. As noted above, more
than a dozen proper-motion surveys of the Pleiades have been made in order to identify cluster members. However, no
single catalog of the cluster has been published that attempts to collect all of those candidate members in a single table and
cross-identify those stars. Another problem is that, while there have been many papers devoted to providing optical
photometry of cluster members, that photometry has been bewilderingly inhomogeneous in terms of the number of
photometric systems used. In § 3 and in the Appendix, we describe our efforts to create a reasonably complete catalog of
candidate Pleiades members and to provide optical photometry transformed to the best of our ability onto a single system.
2. NEW OBSERVATIONAL DATA
2.1. 2MASS "6x" Imaging of the Pleiades
During the final months of Two Micron All Sky Survey (2MASS; Skrutskie et al. 2006) operations, a series of
special observations were carried out that employed exposures 6 times longer than used for the primary survey. These socalled "6x" observations targeted 30 regions of scientific interest including a 3 × 2 area centered on the Pleiades cluster.
The 2MASS 6x data were reduced using an automated processing pipeline similar to that used for the main survey data,
and a calibrated 6x Image Atlas and extracted 6x Point and Extended Source Catalogs (6x-PSC and 6x-XSC) analogous to
the 2MASS All-Sky Atlas, PSC, and XSC have been released as part of the 2MASS Extended Mission. A description of
the content and formats of the 6x image and catalog products, and details about the 6x observations and data reduction, are
given in § A3 of the 2MASS Explanatory Supplement by Cutri et al.4 The 2MASS 6x Atlas and Catalogs may be accessed
via the online services of the NASA/IPAC Infrared Science Archive.5
Figure 1 shows the area on the sky imaged by the 2MASS 6x observations in the Pleiades field. The region was
covered by two rows of scans, each scan being 1° long (in declination) and 8.5 wide in right ascension. Within each row,
the scans overlap by approximately 1 in right
ascension. There are small gaps in coverage in the
declination boundary between the rows, and one
complete scan in the southern row is missing because the
data in that scan did not meet the minimum required
photometric quality. The total area covered by the 6x
Pleiades observations is approximately 5.3 deg2.
Fig. 1 Spatial coverage of the 6 times deeper
(270 "2MASS 6x" observations of the Pleiades. The
kB) 2MASS survey region is approximately centered
on Alcyone, the most massive member of the
Pleiades. The trapezoidal box roughly indicates
the region covered with the shallow IRAC survey
of the cluster core. The star symbols correspond
to the brightest B star members of the cluster. The
red points are the location of objects in the
2MASS 6x Point Source Catalog.
There are approximately 43,000 sources
extracted from the 6x Pleiades observations in the
2MASS 6x-PSC, and nearly 1500 in the 6x-XSC.
Because there are at most about 1000 Pleiades members
expected in this region, only 2% of the 6x-PSC sources are cluster members, and the rest are field stars and background
galaxies. The 6x-XSC objects are virtually all resolved background galaxies. Near-infrared color-magnitude and colorcolor diagrams of the unresolved sources from the 2MASS 6x-PSC and all sources in the 6x-XSC sources from the
Pleiades region are shown in Figures 2 and 3, respectively. The extragalactic sources tend to be redder than most stars, and
the galaxies become relatively more numerous toward fainter magnitudes. Unresolved galaxies dominate the point sources
that are fainter than Ks > 15.5 and redder than J - Ks > 1.2 mag.
Fig. 2 Color-magnitude diagram for
the Pleiades derived from the 2MASS
6x observations. The red dots
correspond to objects identified as
unresolved, whereas the green dots
correspond to extended sources
(primarily background galaxies). The
lack of green dots fainter than K = 16 is
indicative that too few photons are
available to identify sources as
extended—the extragalactic population
presumably increases to fainter
magnitudes.
Fig. 3 Same as Fig. 2, except in this case the
axes are J - H and H - Ks. The extragalactic
objects are very red in both colors.
The 2MASS 6x observations were conducted using the same freeze-frame scanning technique used for the
primary survey (Skrutskie et al. 2006). The longer exposure times were achieved by increasing the "READ2-READ1"
integration to 7.8 s from the 1.3 s used for primary survey. However, the 51 ms "READ1" exposure time was not changed
for the 6x observations. As a result, there is an effective "sensitivity gap" in the 8–11 mag region where objects may be
saturated in the 7.8 s READ2-READ1 6x exposures, but too faint to be detected in the 51 ms READ1 exposures. Because
the sensitivity gap can result in incompleteness and/or flux bias in the photometric overlap regime, the near-infrared
photometry for sources brighter than J = 11 mag in the 6x-PSC was taken from the 2MASS All-Sky PSC during
compilation of the catalog of Pleiades candidate members presented in Table 2 (see § 3).
4
See http://www.ipac.caltech.edu/2mass/releases/allsky/doc/explsup.html.
5
See http://irsa.ipac.caltech.edu.
2.2. Shallow IRAC Imaging
Imaging of the Pleiades with Spitzer was obtained in 2004 April as part of a joint GTO program conducted by the
IRAC instrument team and the Multiband Imaging Photometer for Spitzer (MIPS) instrument team. Initial results of the
MIPS survey of the Pleiades have already been reported in Gorlova et al. (2006). The IRAC observations were obtained as
two astronomical observing requests (AORs). One of them was centered near the cluster center, at R.A. = 03h47m00.0s and
decl. = 24 07 (J2000.0), and consisted of a 12 row by 12 column map, with "frame times" of 0.6 and 12.0 s and two
dithers at each map position. The map steps were 290 in both the column and row direction. The resultant map covers a
region of approximately 1 deg2, and a total integration time per position of 24 s over most of the map. The second AOR
used the same basic mapping parameters, except it was smaller (9 rows by 9 columns) and was instead centered northwest
from the cluster center at R.A. = 03h44m36.0s and decl. = 25 24 . A two-band color image of the AOR covering the center
of the Pleiades is shown in Figure 4. A pictorial guide to the IRAC image providing Greek names for a few of the brightest
stars, and Hertzsprung (1947) numbers for several stars mentioned in § 6 is provided in Figure 5.
Fig. 4 Two-color (4.5 and 8.0 m)
mosaic of the central square degree
of the Pleiades from the IRAC
survey. North is approximately
vertical, and east is approximately to
the left. The bright star nearest the
center is Alcyone; the bright star at
the left of the mosaic is Atlas; and
the bright star at the right of the
mosaic is Electra.
Fig. 5 Finding chart
corresponding
approximately to the region
imaged with IRAC. The large,
five-pointed stars are all of the
Pleiades members brighter than
V = 5.5. The small open circles
correspond to other cluster
members. Several stars with 8
m excesses are labeled by
their HII numbers and are
discussed further in § 6. The
short lines through several of
the stars indicate the size and
position angle of the residual
optical
polarization
(after
subtraction of a constant
foreground component), as
provided in Fig. 6 of Breger
(1986).
We began our analysis with the basic calibrated data (BCDs) from the Spitzer pipeline, using the S13 version of
the Spitzer Science Center pipeline software. Artifact mitigation and masking was done using the IDL tools provided on
the Spitzer contributed software Web site. For each AOR, the artifact-corrected BCDs were combined into single mosaics
for each channel using the post-BCD "MOPEX" package (Makovoz & Marleau 2005). The mosaic images were
constructed with 1.22 × 1.22 pixels (i.e., approximately the same pixel size as the native IRAC arrays).
We derived aperture photometry for stars present in these IRAC mosaics using both APEX (a component of the
MOPEX package) and the "phot" routine in DAOPHOT. In both cases, we used a 3 pixel radius aperture and a sky annulus
from 3 to 7 pixels (except that for channel 4, for the phot package we used a 2 pixel radius aperture and a 2–6 pixel annulus
because that provided more reliable fluxes at low flux levels). We used the flux for zero-magnitude calibrations provided
in the IRAC data handbook (280.9, 179.7, 115.0, and 64.1 Jy for channels 1–4, respectively), and the aperture corrections
provided in the same handbook (multiplicative flux correction factors of 1.124, 1.127, 1.143, and 1.584 for channels 1–4,
inclusive. The channel 4 correction factor is much bigger because it is for an aperture radius of 2 rather than 3 pixels.).
Figures 6 and 7 provide two means to assess the accuracy of the IRAC photometry. The first figure compares the
aperture photometry from APEX to that from phot and shows that the two packages yield very similar results when used in
the same way. For this reason, we have simply averaged the fluxes
from the two packages to obtain our final reported value. The second
figure shows the difference between the derived 3.6 and 4.5 m
magnitudes for Pleiades members. Based on previous studies (e.g.,
Allen et al. 2004), we expected this difference to be essentially zero for
most stars, and the Pleiades data corroborate that expectation. For [3.6]
< 10.5, the rms dispersion of the magnitude difference between the two
channels is 0.024 mag. Assuming that each channel has similar
uncertainties, this indicates an internal 1 accuracy of order 0.017
mag. The absolute calibration uncertainty for the IRAC fluxes is
currently estimated at of order 0.02 mag. Figure 7 also shows that
fainter than [3.6] = 10.5 (spectral type later than about M0), the [3.6] [4.5] color for M dwarfs departs slightly from zero, becoming
increasingly redder to the limit of the data (about M6).
Fig. 6 Comparison of aperture photometry for Pleiades members
derived
from the IRAC 3.6 m mosaic using the Spitzer APEX package and
the IRAF implementation of DAOPHOT.
Fig. 7 Difference between aperture photometry for Pleiades
members for IRAC channels 1 and 2. The [3.6] - [4.5] color
begins to depart from essentially zero at magnitudes of 10.5,
corresponding approximately to spectral type M0 in the
Pleiades.
3. A CATALOG OF PLEIADES CANDIDATE MEMBERS
If one limits oneself to only stars visible with the naked eye, it is easy to identify which stars are members of the
Pleiades—all of the stars within a degree of the cluster center that have V < 6 are indeed members. However, if one were to
try to identify the M dwarf stellar members of the cluster (roughly 14 < V < 23), only of order 1% of the stars toward the
cluster center are likely to be members, and it is much harder to construct an uncontaminated catalog. The problem is
exacerbated by the fact that the Pleiades is old enough that mass segregation through dynamical processes has occurred,
and therefore one has to survey a much larger region of the sky in order to include all of the M dwarf members.
The other primary difficulty in constructing a comprehensive member catalog for the Pleiades is that the pedigree
of the candidates varies greatly. For the best-studied stars, astrometric positions can be measured over temporal baselines
ranging up to a century or more, and the separation of cluster members from field stars in a vector point diagram (VPD)
can be extremely good. In addition, accurate radial velocities and other spectral indicators are available for essentially all
of the bright cluster members, and these further allow membership assessment to be essentially definitive. Conversely, at
the faint end (for stars near the hydrogen-burning mass limit in the Pleiades), members are near the detection limit of the
existing wide-field photographic plates, and the errors on the proper motions become correspondingly large, causing the
separation of cluster members from field stars in the VPD to become poor. These stars are also sufficiently faint that
spectra capable of discriminating members from field dwarfs can only be obtained with 8 m class telescopes, and only a
very small fraction of the faint candidates have had such spectra obtained. Therefore, any comprehensive catalog created
for the Pleiades will necessarily have stars ranging from certain members to candidates for which very little is known and
where the fraction of spurious candidate members increases to lower masses.
In order to address the membership uncertainties and biases, we have chosen a sliding scale for inclusion in our
catalog. For all stars, we require that the available photometry yields location in color-color and color-magnitude diagrams
consistent with cluster membership. For the stars with well-calibrated photoelectric photometry, this means the star should
not fall below the Pleiades single-star locus by more than about 0.2 mag or above that locus by more than about 1.0 mag
(the expected displacement for a hierarchical triple with three nearly equal mass components). For stars with only
photographic optical photometry, where the 1 uncertainties are of order 0.1–0.2 mag, we still require the star's
photometry to be consistent with membership, but the allowed displacements from the single-star locus are considerably
larger. Where accurate radial velocities are known, we require that the star be considered a radial velocity member based
on the paper where the radial velocities were presented. Where stars have been previously identified as nonmembers based
on photometric or spectroscopic indices, we adopt those conclusions.
Two other relevant pieces of information are sometimes available. In some cases, individual proper-motion
membership probabilities are provided by the various membership surveys. If no other information is available, and if the
membership probability for a given candidate is less than 0.1, we exclude that star from our final catalog. However, often a
star appears in several catalogs; if it appears in two or more proper-motion membership lists, we include it in the final
catalog even if P < 0.1 in one of those catalogs. Second, an entirely different means to identify candidate Pleiades
members is via flare star surveys toward the cluster (Haro et al. 1982; Jones 1981). A star with a formally low membership
probability in one catalog but whose photometry is consistent with membership and that was identified as a flare star is
retained in our catalog.
Further details of the catalog construction are provided in the Appendix, as are details of the means by which the
B, V, and I photometry have been homogenized. A full discussion and listing of all of the papers from which we have
extracted astrometric and photometric information is also provided in the Appendix. Here we simply provide a very brief
description of the inputs to the catalog.
We include candidate cluster members from the following proper-motion surveys: Trumpler (1921), Hertzsprung
(1947), Jones (1981), Pels & Lub (as reported in van Leeuwen et al. 1986), Stauffer et al. (1991), Artyukhina (1969),
Hambly et al. (1993), Pinfield et al. (2000), Adams et al. (2001), and Deacon & Hambly (2004). Another important
compilation that provides the initial identification of a significant number of low-mass cluster members is the flare star
catalog of Haro et al. (1982). Table 1 provides a brief synopsis of the characteristics of the candidate member catalogs
from these papers. The Trumpler paper is listed twice in Table 1 because there are two membership surveys included in
that paper, with differing spatial coverages and different limiting magnitudes.
Table 1 CITED IN TEXT | ASCII | TYPESET IMAGE
Table
1 Pleiades
Membership Surveys
Used as Sources
Go to: Table 2
Pleiades Membership Surveys Used as Sources
Reference
Area
Covered
(deg2)
Magnitude
Range
(and Band)
Number of
Candidates
Name
Prefix
Trumpler (1921)...
3
2.5 < B < 14.5
174
Tr
Trumpler (1921)a...
24
2.5 < B < 10
72
Tr
Hertzsprung (1947)...
4
2.5 < V < 15.5
247
HII
Artyukhina (1969)...
60
2.5 < B < 12.5
Haro et al. (1982)...
20
11 < V < 17.5
519
HCG
van Leeuwen et al.
(1986)...
80
2.5 < B < 13
193
Pels
Stauffer et al. (1991)...
16
14 < V < 18
225
SK
Hambly et al. (1993)...
23
10 < I < 17.5
440
HHJ
Pinfield et al. (2000)...
6
13.5 < I < 19.5
339
BPL
Adams et al. (2001)...
300
8 < Ks < 14.5
1200
...
Deacon & Hambly
(2004)...
75
10 < R < 19
916
DH
a
200
AK
The Trumpler paper is listed twice because there are two membership surveys
included in that paper, with differing spatial coverages and different limiting
magnitudes.
In our final catalog, we have attempted to follow the standard naming convention whereby the primary name is
derived from the paper where it was first identified as a cluster member. An exception to this arises for stars with both
Trumpler (1921) and Hertzsprung (1947) names, where we use the Hertzsprung numbers as the standard name because
that is the most commonly used designation for these stars in the literature. The failure for the Trumpler numbers to be
given precedence in the literature perhaps stems from the fact that the Trumpler catalog was published in the Lick
Observatory Bulletins as opposed to a refereed journal. In addition to providing a primary name for each star, we provide
cross-identifications to some of the other catalogs, particularly where there is existing photometry or spectroscopy of that
star using the alternate names. For the brightest cluster members, we provide additional cross-references (e.g., Greek
names, Flamsteed numbers, HD numbers).
For each star, we attempt to include an estimate for Johnson B and V, and for Cousins I (IC). Only a very small
fraction of the cluster members have photoelectric photometry in these systems, unfortunately. Photometry for many of the
stars has often been obtained in other systems, including Walraven, Geneva, Kron, and Johnson. We have used previously
published transformations from the appropriate indices in those systems to Johnson BV or Cousins I. In other cases,
photometry is available in a natural I-band system, primarily for some of the relatively faint cluster members. We have
attempted to transform those I-band data to IC by deriving our own conversion using stars for which we already have an IC
estimate as well as the natural I measurement. Details of these issues are provided in the Appendix.
Finally, we have cross-correlated the cluster candidates catalog with the 2MASS All-Sky PSC and also with the
6x-PSC for the Pleiades. For every star in the catalog, we obtain JHKs photometry and 2MASS positions. Where we have
both main survey 2MASS data and data from the 6x catalog, we adopt the 6x data for stars with J > 11, and data from the
standard 2MASS catalog otherwise. We verified that the two catalogs do not have any obvious photometric or astrometric
offsets relative to each other. The coordinates we list in our catalog are entirely from these 2MASS sources, and hence
they inherit the very good and homogeneous 2MASS positional accuracies of order 0.1 rms.
We have then plotted the candidate Pleiades members in a variety of color-magnitude diagrams and color-color
diagrams and required that a star must have photometry that is consistent with cluster membership. Figure 8 illustrates this
process and indicates why (for example) we have excluded HII 1695 from our final catalog.
Fig. 8 Ks vs. Ks - [4.5] CMD for
Pleiades
candidate
members,
illustrating why we have excluded HII
1695 from the final catalog of cluster
members. The "X" symbol marks the
location of HII 1695 in this diagram.
Table 2 provides the collected data for the 1417 stars we have retained as candidate Pleiades members. The first two
columns are the J2000.0 right ascension and declination from 2MASS; the next are the 2MASS JHKs photometry and their
uncertainties, and the 2MASS photometric quality flag ("ph-qual"). If the number following the 2MASS quality flag is a 1,
the 2MASS data come from the 2MASS All-Sky PSC; if it is a 2, the data come from the 6x-PSC. The next three columns
provide the B, V, and IC photometry, followed by a flag that indicates the provenance of that photometry. The last column
provides the most commonly used names for these stars. The hydrogen-burning mass limit for the Pleiades occurs at about
V = 22, I = 18, Ks = 14.4. Fifty-three of the candidate members in the catalog are fainter than this limit and hence should be
substellar if they are indeed Pleiades members.
R.A.
(J2000.0)
(deg)
Dec.
(J2000.0)
(deg)
51.898273.
..
Referenc
e
Name
s
11.8
5
22
DH
001
10.3
0
...
4
Pels
121
...
...
14.3
4
22
DH
003
AAA
1
...
...
15.0
5
22
DH
004
9.723
±
0.016
AAA
1
12.5
9
11.7
5
...
9
AKIII
59
12.44
2±
0.029
12.15
3±
0.015
AAA
1
...
...
14.7
5
22
DH
006
10.42
9±
0.019
10.00
6±
0.029
9.924
±
0.016
AAA
1
...
...
11.1
6
22
DH
007
25.65230
4
8.459
±
0.015
8.314
±
0.059
8.270
±
0.031
AAA
1
9.90
9.43
...
9
AKIII
79
52.409874.
..
24.51054
6
10.31
8±
0.023
9.856
±
0.028
9.698
±
0.016
AAA
1
...
...
11.1
9
22
DH
008
52.494766.
..
23.37185
9
12.79
9±
0.018
12.20
6±
0.019
11.94
6±
0.018
AAA
1
...
...
14.3
9
22
DH
009
52.534420.
..
22.64416
3
13.68
0±
12.96
8±
12.77
0±
AAA
1
...
...
15.1
6
22
DH
010
J
H
Ks
24.52866
0
10.78
1±
0.025
10.06
6±
0.030
9.892
±
0.017
51.925262.
..
23.80368
8
9.066
±
0.013
8.754
±
0.009
51.976067.
..
24.93647
8
12.88
0±
0.019
52.006481.
..
23.07849
9
52.168613.
..
phquala
B
V
IC
AAA
1
...
...
8.679
±
0.014
AAA
1
10.9
6
12.21
9±
0.030
11.98
1±
0.016
AAA
1
13.52
5±
0.022
12.91
9±
0.022
12.61
9±
0.021
25.60778
2
10.19
8±
0.019
9.883
±
0.029
52.200249.
..
25.57584
2
13.07
2±
0.019
52.203186.
..
26.49935
0
52.355843.
..
Table
2 Pleiade
s
Members:
Literature
Photometr
y
0.022
0.024
0.023
52.639614.
..
26.21576
7
11.00
7±
0.018
10.40
0±
0.027
10.27
9±
0.016
AAA
1
...
...
11.8
6
22
DH
011
52.647411.
..
23.05233
4
15.31
9±
0.051
14.56
5±
0.060
14.26
0±
0.073
AAA
1
...
...
16.6
8
22
DH
012
52.656086.
..
26.34610
0
14.23
9±
0.026
13.52
4±
0.037
13.29
9±
0.032
AAA
1
...
...
15.7
3
22
DH
013
52.799591.
..
25.16510
0
11.48
9±
0.016
10.78
4±
0.020
10.60
2±
0.014
AAA
1
...
...
12.6
5
22
DH
014
52.810955.
..
25.98114
8
11.99
1±
0.018
11.29
1±
0.022
11.09
1±
0.018
AAA
1
...
...
13.3
0
22
DH
015
52.873249.
..
26.50352
5
13.61
3±
0.022
13.01
8±
0.026
12.75
8±
0.022
AAA
1
...
...
15.2
8
22
DH
016
52.890076.
..
26.26550
7
9.514
±
0.016
9.222
±
0.021
9.068
±
0.015
AAA
1
11.4
5
10.7
7
...
4
Pels
008
53.001957.
..
23.77490
0
11.32
9±
0.017
10.68
6±
0.019
10.52
0±
0.016
AAA
1
15.2
7
13.9
5
...
4
Pels
109
53.032749.
..
23.23265
5
13.13
5±
0.019
12.48
6±
0.019
12.25
4±
0.018
AAA
1
...
...
14.7
1
22
DH
017
Note.— Table 2 is available in its entirety via the link to the machine-readable version above.
a
Standard 2MASS photometric data quality flag for JHKs, in that order. If the number following the
2MASS quality flags is a 1, the 2MASS data come from the standard 2MASS catalog; if it is a 2, the data
come from the deep catalog.
Table 3 provides the IRAC [3.6], [4.5], [5.8], and [8.0] photometry we have derived for Pleiades candidate members
included within the region covered by the IRAC shallow survey of the Pleiades (see § 2). The brightest stars are saturated
even in our short integration frame data, particularly for the more sensitive 3.6 and 4.5 m channels. At the faint end, we
provide photometry only for 3.6 and 4.5 m because the objects are undetected in the two longer wavelength channels. At
the "top" and "bottom" of the survey region, we have incomplete wavelength coverage for a band of width about 5 , and
for stars in those areas we report only photometry in either the 3.6 and 5.8 bands or the 4.5 and 8.0 bands.
Table 3 Pleiades Members: IRAC Photometry
Name
[3.6]
[4.5]
[5.8]
[8]
HHJ 107... 12.550 12.514 12.474 12.388
HCG 96...
11.869
11.881
11.805
11.824
DH 257...
9.604
9.608
9.604
9.554
SK 646...
11.318
11.273
11.204
11.215
HII 97...
...
9.760
...
9.666
Pels 056...
9.188
9.214
9.164
9.165
HCG 112...
11.711
11.646
11.623
11.620
SK 622...
11.686
11.656
11.699
11.575
HCG 115...
11.450
11.434
11.316
11.437
HII 153...
7.163
7.205
7.183
7.198
HII 174...
...
9.325
...
9.285
HII 173...
8.798
8.812
8.763
8.768
HCG 125...
11.641
11.594
11.564
11.540
Pels 043...
9.673
...
9.673
...
SK 609...
15.572
15.663
15.857
15.685
AK 1B146...
8.189
8.175
8.153
8.166
HHJ 218...
12.309
12.285
12.268
12.494
HCG 126...
11.220
11.205
11.181
11.193
SK 596...
11.141
11.112
11.039
11.061
HCG 129...
...
11.276
...
11.282
HCG 134...
11.111
11.071
11.032
11.033
HCG 131...
9.941
9.982
9.980
9.911
HII 250...
...
9.083
...
9.023
Pels 059...
9.950
9.991
9.951
9.934
HHJ 235...
12.167
12.085
11.940
12.080
HCG 138...
11.200
11.132
11.096
11.124
HCG 143...
11.355
11.303
11.258
11.291
HHJ 100...
12.588
12.551
12.462
12.459
HCG 152...
10.858
10.874
10.857
10.846
HHJ 68...
13.007
12.914
13.089
12.688
HCG 157...
12.194
...
12.045
...
HII 380...
10.169
10.192
10.173
10.161
HHJ 46...
13.149
13.085
13.170
12.898
Pels 041...
9.740
9.774
9.674
9.698
HII 430...
9.509
9.459
9.356
9.465
HHJ 24...
13.410
13.345
13.298
13.784
HCG 166...
12.168
12.111
11.908
12.291
HII 447...
...
...
5.522
5.528
HII 468...
...
...
4.010
3.910
HHJ 183...
12.458
12.426
12.431
12.207
HII 489...
8.857
8.885
8.864
8.814
DH 367...
12.770
12.717
12.594
12.589
HHJ 139...
...
12.495
...
12.628
HII 514...
9.019
9.006
8.955
8.978
SK 534...
11.269
11.262
11.163
11.241
HII 531...
7.715
7.731
7.703
7.719
HHJ 164...
12.469
12.381
12.354
12.410
HII 554...
...
10.456
...
10.427
HII 563...
...
...
4.620
4.580
HHJ 14...
13.479
13.461
13.624
13.430
HCG 180...
12.854
12.775
12.763
12.697
HII 559...
...
10.120
...
10.097
HII 566...
...
10.722
...
10.681
HII 571...
9.175
9.151
9.161
9.097
SK 526...
...
10.744
...
10.698
HCG 181...
...
11.173
...
11.141
HCG 178...
10.938
10.960
10.774
10.767
HII 590...
...
10.552
...
10.503
HII 625...
9.317
9.330
9.323
9.250
HHJ 273...
...
12.014
...
11.962
DH 392...
11.972
11.934
11.827
12.053
HII 652...
7.426
7.455
7.422
7.278
HHJ 99...
13.052
13.021
13.180
13.126
HHJ 106...
12.975
12.923
12.953
12.619
HII 676...
10.149
10.131
10.105
10.132
HII 673...
10.938
10.944
10.910
10.997
HHJ -293...
12.227
...
12.060
...
HII 686...
10.100
10.116
10.126
10.084
HII 697...
7.703
7.661
7.656
7.597
HII 708...
8.547
8.507
8.472
8.497
HII 717...
6.538
6.600
6.592
6.611
HCG 196...
10.767
10.810
10.767
10.749
HHJ 130...
12.724
12.695
12.618
12.781
HII 738...
8.819
8.845
8.773
8.762
HII 745...
8.002
8.007
7.975
7.988
HCG 195...
10.648
10.611
10.561
10.587
HII 746...
9.300
9.328
9.292
9.280
HII 740...
...
10.495
...
10.425
HII 762...
10.556
10.571
10.533
10.479
HII 761...
8.736
8.745
8.722
8.694
HII 793...
10.529
10.621
10.493
10.431
DH 403...
14.486
14.432
14.604
13.951
BPL 77...
11.528
11.500
11.390
55.505
HII 785...
...
...
4.050
4.030
HII 799...
10.140
10.162
10.107
10.178
BPL 79...
14.069
14.009
13.914
13.960
HII 804...
7.345
7.323
7.340
7.338
HHJ 166...
12.469
12.405
12.364
12.410
BPL 81...
14.498
14.477
14.190
14.098
HII 813...
10.355
10.340
10.283
10.288
HII 817...
...
9.900
...
5.737
BPL 82...
11.539
11.510
11.449
11.582
SK 497...
11.198
11.172
11.112
11.107
HHJ 27...
13.407
13.418
13.113
13.243
DH 412...
13.309
13.297
13.161
13.409
HHJ 127...
12.785
12.736
12.774
12.486
SK 491...
11.042
11.052
10.994
11.016
HII 870...
9.133
9.134
9.109
9.082
HII 859...
...
6.454
...
6.422
10.656
10.684
10.635
SK 490...
10.672
HHJ 363...
11.628
11.536
11.505
11.552
BPL 88...
12.871
12.811
12.703
12.560
SK 488...
...
11.196
...
11.137
HHJ 194...
12.575
12.635
12.546
12.712
HII 879...
...
10.081
...
10.066
HHJ 435...
10.767
10.761
10.733
10.702
HII 883...
...
10.195
...
10.125
HII 890...
10.727
10.715
10.649
10.696
HII 916...
...
9.524
...
9.479
HII 930...
10.459
10.472
10.424
10.459
HHJ 56...
12.475
12.422
12.572
12.562
HII 956...
7.092
7.131
7.079
7.069
HII 980...
...
...
...
4.190
HHJ 105...
...
12.724
...
12.609
DH 441...
11.766
11.693
11.531
11.636
HII 996...
...
8.932
...
8.878
HCG 218...
12.534
12.462
12.443
12.465
HHJ 249...
12.259
12.254
12.173
12.154
HCG 219...
10.889
10.851
10.794
10.801
HHJ 326...
11.786
11.731
11.618
11.719
HII 1028...
7.078
7.120
7.113
7.085
HII 1015...
...
9.016
...
8.965
HHJ 161...
12.285
12.181
12.275
12.181
HII 1039...
9.798
...
9.764
...
HII 1032...
9.143
9.129
9.144
9.071
HII 1061...
10.298
10.323
10.245
10.240
HII 1084...
7.052
...
7.043
...
HHJ 140...
12.439
12.385
12.431
12.409
HII 1094...
10.549
10.546
10.479
10.613
HII 1100...
9.285
9.321
9.302
9.264
HII 1117...
8.497
8.524
8.543
8.482
HII 1110...
10.227
10.284
10.240
10.208
HII 1122...
8.149
8.146
8.174
8.132
HII 1124...
9.858
9.843
9.847
9.785
HHJ 104...
12.705
...
12.752
...
DH 467...
11.379
11.297
11.253
11.271
HII 1173...
...
10.855
...
10.781
HHJ 247...
11.836
...
11.660
...
HCG 244...
10.955
10.921
10.873
10.870
HII 1215...
8.997
...
8.953
...
HHJ 257...
11.948
11.924
11.906
11.714
HHJ 174...
12.506
...
...
...
HHJ 299...
11.797
11.744
11.694
11.663
HII 1234...
6.729
6.743
6.712
6.679
HHJ 252...
11.927
11.863
11.740
11.836
HII 1280...
10.548
10.602
10.528
10.575
HII 1286...
10.378
...
10.289
...
HII 1284...
7.617
7.627
7.571
7.589
HII 1298...
9.778
9.784
9.800
9.741
HCG 253...
11.625
11.513
11.511
11.559
HII 1306...
9.798
9.811
9.773
9.747
HHJ 37...
13.241
13.158
13.051
13.073
HII 1321...
10.438
10.413
10.361
10.372
HII 1309...
8.270
8.285
8.244
8.259
HHJ 92...
12.816
12.717
12.684
12.753
HII 1332...
9.969
10.016
9.991
9.958
HCG 258...
10.839
10.766
10.673
10.750
HII 1338...
7.463
7.483
7.503
7.476
HII 1348...
9.622
9.651
9.622
9.575
HII 1355...
10.016
10.024
9.991
10.035
HII 1362...
7.637
7.661
7.641
7.618
HII 1380...
6.914
6.952
6.961
6.962
HII 1375...
...
6.323
6.311
6.318
HCG 266...
12.175
12.115
12.078
11.993
HII 1384...
...
6.984
...
6.979
HII 1397...
7.181
7.196
7.192
7.184
HHJ 198...
12.717
12.668
12.541
12.431
HCG 269...
11.965
...
11.859
...
HII 1425...
7.342
...
7.323
...
HII 1431...
6.640
6.652
6.644
6.674
HCG 273...
11.175
11.142
11.093
11.240
HCG 277...
10.658
...
10.553
...
HII 1454...
...
10.071
...
10.026
DH 523...
12.495
12.493
12.472
12.348
HII 1516...
10.219
10.252
10.211
10.217
HII 1514...
8.940
8.936
8.910
8.913
HII 1532...
10.510
10.472
10.450
10.403
HII 1531...
10.260
10.275
10.191
10.195
HHJ 26...
13.951
13.840
13.870
13.714
HHJ 152...
12.384
...
12.228
...
HHJ 438...
11.138
11.076
11.014
11.048
HCG 295...
11.248
11.254
11.248
11.190
HHJ 122...
12.698
...
12.635
...
HII 1613...
8.562
8.565
8.514
8.535
HHJ 240...
12.232
12.184
12.131
12.102
HII 1726...
7.905
7.904
7.883
7.891
HCG 307...
11.929
11.899
11.830
11.887
HHJ 156...
12.458
...
12.384
...
HHJ 225...
12.370
12.237
12.267
12.356
DH 555...
13.164
13.108
12.971
13.066
HCG 311...
12.095
12.014
11.969
12.067
HII 1762...
7.306
7.312
7.342
7.332
HCG 315...
11.888
...
11.717
...
HHJ 336...
11.528
...
11.462
...
HII 1797...
8.729
...
8.662
...
HII 1794...
8.871
8.887
8.812
8.797
HII 1785...
10.569
10.594
10.552
10.563
HHJ 188...
12.467
12.384
12.441
12.456
HII 1827...
10.298
10.272
10.163
10.217
HII 1856...
8.648
8.637
8.613
8.626
HCG 328...
12.677
12.611
12.659
12.887
HII 1876...
6.575
6.578
6.595
6.592
HCG 324...
11.169
11.146
11.114
11.146
HCG 327...
12.521
12.483
12.455
12.506
HHJ 184...
12.586
12.555
12.573
12.553
HII 1912...
7.798
7.834
7.783
7.790
HCG 335...
12.866
12.826
12.760
12.820
HCG 337...
11.627
11.558
11.584
11.665
HHJ 44...
13.179
13.093
13.023
12.903
HHJ 207...
12.389
12.304
12.221
12.330
HII 2027...
8.788
8.836
8.774
8.784
HII 2034...
9.878
9.922
9.923
9.865
DH 593...
14.247
14.270
14.972
14.808
HHJ 8...
13.674
13.656
13.473
13.729
HHJ 231...
12.236
12.194
12.171
12.189
HCG 354...
10.566
...
10.486
...
HII 2147...
8.558
8.615
8.514
8.549
HII 2168...
...
...
3.840
3.820
HII 2195...
7.600
7.622
7.601
7.588
DH 610...
13.173
...
13.132
...
HII 2284...
9.366
9.384
9.353
9.319
HII 2311...
9.445
...
9.395
...
HHJ 142...
12.426
...
12.224
...
Because Table 2 is an amalgam of many previous catalogs, each of which have different spatial coverage, magnitude
limits, and other idiosyncrasies, it is necessarily incomplete and inhomogeneous. It also certainly includes some
nonmembers. For V < 12, we expect very few nonmembers because of the extensive spectroscopic data available for those
stars; the fraction of nonmembers will likely increase to fainter magnitudes, particularly for stars located far from the
cluster center. The catalog is simply an attempt to collect all of the available data, identify some of the nonmembers, and
eliminate duplications. We hope that it will also serve as a starting point for future efforts to produce a "cleaner" catalog.
Figure 9 shows the distribution on the sky of the stars in Table 2. The complete spatial distribution of all members of
the Pleiades may differ slightly from what is shown due to the inhomogeneous properties of the proper-motion surveys.
However, we believe that those effects are relatively small and the distribution shown is mostly representative of the
parent population. One thing that is evident in Figure 9 is mass segregation—the highest mass cluster members are much
more centrally located than the lowest mass cluster members. This fact is reinforced by calculating the cumulative number
of stars as a function of distance from the cluster center for different absolute magnitude bins. Figure 10 illustrates this
fact. Another property of the Pleiades illustrated by Figure 10 is that the cluster appears to be elongated parallel to the
Galactic plane, as expected from n-body simulations of galactic clusters (Terlevich 1987). Similar plots showing the
flattening of the cluster and evidence for mass segregation for the V < 12 cluster members were provided by Raboud &
Mermilliod (1998).
Fig. 9 Spatial plot of the
candidate Pleiades members
from Table 2. The large star
symbols are members brighter
than Ks = 6; the open circles are
stars with 6 < Ks < 9; and the
dots are candidate members
fainter than Ks = 9. The solid
line is parallel to the Galactic
plane.
Fig.
10 Cumul
ative
radial
density
profiles for
Pleiades
members in
several
magnitude
ranges:
heavy,
longdashed
line, Ks <
6; dots, 6
< Ks < 9;
shortdashed
line, 9 < Ks
< 12; light,
longdashed
line, Ks >
12.
4. EMPIRICAL PLEIADES ISOCHRONES
AND COMPARISON TO MODEL ISOCHRONES
Young, nearby, rich open clusters like the Pleiades
can and should be used to provide template data that can help interpret observations of more distant clusters or to test
theoretical models. The identification of candidate members of distant open clusters is often based on plots of stars in a
color-magnitude diagram, overlaid on which is a line meant to define the single-star locus at the distance of the cluster.
The stars lying near or slightly above the locus are chosen as possible or probable cluster members. The data we have
collected for the Pleiades provide a means to define the single-star locus for 100 Myr, solar metallicity stars in a variety of
widely used color systems down to and slightly below the hydrogen-burning mass limit. Figures 11 and 12 illustrate the
appearance of the Pleiades stars in two of these diagrams, and the single-star locus we have defined. The curve defining
the single-star locus was drawn entirely "by eye." It is displaced slightly above the lower envelope to the locus of stars to
account for photometric uncertainties (which increase to fainter magnitudes). We attempted to use all of the information
available to us, however. That is, there should also be an upper envelope to the Pleiades locus in these diagrams, since
equal-mass binaries should be displaced above the single-star sequence by 0.7 mag (and one expects very few systems of
higher multiplicity). Therefore, the single-star locus was defined with that upper envelope in mind. Table 4 provides the
single-star loci for the Pleiades for BVICJKs plus the four IRAC channels. We have dereddened the empirical loci by the
canonical mean extinction to the Pleiades of AV = 0.12 (and, correspondingly, AB = 0.16, AI = 0.07, AJ = 0.03, and AK =
0.01, as per the reddening law of Rieke & Lebofsky 1985).
Fig. 11 V vs. (V - I)c CMD for Pleiades members with photoelectric photometry. The solid curve is the "by-eye" fit to
the single-star locus for Pleiades members.
Fig. 12 Ks vs. Ks - [3.6] CMD for Pleiades candidate
members from Table 2 (dots) and from deeper imaging of a
set of Pleiades VLM and brown dwarf candidate members
from P. Lowrance et al. (2007, in preparation) (squares).
The solid curve is the single-star locus from Table 4.
(51 kB)
B
V
IC
Ks
[3.6] [4.5] [5.8]
[8]
6.598...
6.600
6.574
6.592
6.602
6.615
6.602
6.602
6.706...
6.700
6.665
6.671
6.682
6.695
6.682
6.682
6.814...
6.800
6.755
6.750
6.761
6.775
6.762
6.762
6.922...
6.900
6.848
6.834
6.841
6.855
6.841
6.841
7.030...
7.000
6.940
6.910
6.920
6.935
6.921
6.921
7.142...
7.100
7.030
6.982
6.990
7.005
6.991
6.991
7.254...
7.200
7.115
7.039
7.050
7.064
7.050
7.049
7.370...
7.300
7.200
7.104
7.109
7.124
7.108
7.107
7.490...
7.400
7.283
7.162
7.168
7.183
7.165
7.164
7.610...
7.500
7.367
7.228
7.238
7.252
7.233
7.231
7.730...
7.600
7.450
7.287
7.297
7.311
7.291
7.288
7.850...
7.700
7.533
7.345
7.347
7.360
7.339
7.336
7.968...
7.800
7.615
7.387
7.396
7.409
7.387
7.384
8.084...
7.900
7.698
7.428
7.436
7.449
7.426
7.422
8.200...
8.000
7.780
7.469
7.475
7.488
7.465
7.460
Table 4 Single-Star Pleiades Loci
8.320...
8.100
7.840
7.508
7.515
7.527
7.503
7.498
8.440...
8.200
7.900
7.546
7.555
7.567
7.542
7.536
8.564...
8.300
7.975
7.591
7.594
7.606
7.580
7.575
8.692...
8.400
8.050
7.648
7.654
7.665
7.638
7.632
8.820...
8.500
8.125
7.701
7.703
7.715
7.687
7.680
8.936...
8.600
8.200
7.762
7.762
7.774
7.744
7.737
The other benefit to constructing the new catalog is that it can provide an improved comparison data set to test
theoretical isochrones. The new catalog provides homogeneous photometry in many photometric bands for stars ranging
from several solar masses down to below 0.1 M . We take the distance to the Pleiades as 133 pc and refer the reader to
Soderblom et al. (2005) for a discussion and a listing of the most recent determinations. The age of the Pleiades is not as
well-defined but is probably somewhere between 100 and 125 Myr (Meynet et al. 1993; Stauffer et al. 1999). We adopt
100 Myr for the purposes of this discussion; our conclusions relative to the theoretical isochrones would not be affected
significantly if we instead chose 125 Myr. As noted above, we adopt AV = 0.12 as the mean Pleiades extinction and apply
that value to the theoretical isochrones. A small number of Pleiades members have significantly larger extinctions (Breger
1986; Stauffer & Hartmann 1987), and we have dereddened those stars individually to the mean cluster reddening.
Figures 13 and 14 compare theoretical 100 Myr isochrones from Siess et al. (2000) and Baraffe et al. (1998) to the
Pleiades member photometry from Table 2 for stars for which we have photoelectric photometry. Neither set of isochrones
are a good fit to the V - I based color-magnitude diagram. For Baraffe et al. (1998) this is not a surprise because they
illustrated that their isochrones are too blue in V - I for cool stars in their paper and ascribed the problem as likely the result
of an incomplete line list, resulting in too little absorption in the V band. For Siess et al. (2000) the poor fit in the V - I
CMD is somewhat unexpected in that they transform from the theoretical to the observational plane using empirical colortemperature relations. In any event, it is clear that neither model isochrones match the shape of the Pleiades locus in the V
versus V - I plane, and therefore use of these V - I based isochrones for younger clusters is not likely to yield accurate
results (unless the color-Teff relation is recalibrated, as described, e.g., in Jeffries & Oliveira 2005). On the other hand, the
Baraffe et al. (1998) model provides a quite good fit to the Pleiades single-star locus for an age of 100 Myr in the K versus
I - K plane.6 This perhaps lends support to the hypothesis that the misfit in the V versus V - I plane is due to missing opacity
in their V-band atmospheres for low-mass stars (see also Chabrier et al. 2000 for further evidence in support of this idea).
The Siess et al. (2000) isochrones do not fit the Pleiades locus in the K versus I - K plane particularly well, being too faint
near I - K = 2 and too bright for I - K > 2.5.
Fig. 13 V vs. (V I)c
CMD
for
Pleiades candidate
members
from
Table 2 for which
we
have
photoelectric
photometry,
compared
to
theoretical
isochrones from
Siess et al. (2000)
(left) and from
Baraffe et
al.
(1998) (right). For
the left panel, the
curves correspond
to 10, 50, and 100
Myr and a ZAMS;
the right panel
includes curves for
50 and 100 Myr
and a ZAMS.
(74 kB)
Fig. 14 K vs. (I - K)
CMD for Pleiades
candidate members
from
Table
2,
compared
to
theoretical
isochrones from Siess
et al. (2000) (left) and
from Baraffe et al.
(1998) (right). The
curves correspond to
50 and 100 Myr and a
ZAMS.
6
These isochrones are calculated for the standard K filter, rather than Ks. However, the difference in location of
the isochrones in these plots because of this should be very slight, and we do not believe our conclusions are significantly
affected.
5. IDENTIFICATION OF NEW VERY LOW-MASS CANDIDATE MEMBERS
The highest spatial density for Pleiades members of any mass should be at the cluster center. However, searches
for substellar members of the Pleiades have generally avoided the cluster center because of the deleterious effects of
scattered light from the high-mass cluster members and because of the variable background from the Pleiades reflection
nebulae. The deep 2MASS and IRAC 3.6 m imaging and 4.5 m imaging provide accurate photometry to well below
the hydrogen-burning mass limit and are less affected by the nebular emission than shorter wavelength images. We
therefore expect that it should be possible to identify a new set of candidate Pleiades substellar members by combining our
new near- and mid-infrared photometry.
The substellar mass limit in the Pleiades occurs at about Ks = 14.4, near the limit of the 2MASS All-Sky PSC. As
illustrated in Figure 15, the deep 2MASS survey of the Pleiades should easily detect objects at least 2 mag fainter than the
substellar limit. The key to actually identifying those objects and separating them from the background sources is to find
color-magnitude or color-color diagrams that separate the Pleiades members from the other objects. As shown in Figure
15, late-type Pleiades members separate fairly well from most field stars toward the Pleiades in a Ks versus Ks - [3.6] colormagnitude diagram. However, as illustrated in Figure 2, in the Ks magnitude range of interest there is also a large
population of red galaxies, and they are in fact the primary contaminants to identifying Pleiades substellar objects in the Ks
versus Ks - [3.6] plane. Fortunately, most of the contaminant galaxies are slightly resolved in the 2MASS and IRAC
imaging, and we have found that we can eliminate most of the red galaxies by their nonstellar image shape.
Figure 15 shows the first step in our process of identifying new very low-mass members of the Pleiades. The red
plus symbols are the known Pleiades members from Table 2. The red open circles are candidate Pleiades substellar
members from deep imaging surveys published in the literature, mostly of parts of the cluster exterior to the central square
degree, where the IRAC photometry is from P. Lowrance et al. (2007, in preparation). The blue, filled circles are field M
and L dwarfs, placed at the distance of the Pleiades, using photometry from Patten et al. (2006). Because the Pleiades is
100 Myr, its very low-mass stellar and substellar objects will be displaced about 0.7 mag above the locus of the field M
and L dwarfs according to the Baraffe et al. (1998) and Chabrier et al. (2000) models, in accord with the location in the
diagram of the previously identified, candidate VLM and substellar objects. The trapezoidal shaped region outlined with a
dashed line is the region in the diagram that we define as containing candidate new VLM and substellar members of the
Pleiades. We place the faint limit of this region at Ks = 16.2 in order to avoid the large apparent increase in faint, red
objects for Ks > 16.2, caused largely by increasing errors in the Ks photometry. Also, the 2MASS extended object flags
cease to be useful fainter than about Ks = 16.
Fig. 15 Ks vs. Ks - [3.6] CMD for the
objects in the central 1 deg2 of the
Pleiades, combining data from the
IRAC shallow survey and 2MASS. The
symbols are defined within the figure
(and see text for details). The dashedline box indicates the region within
which we have searched for new
candidate Pleiades
VLM
and
substellar members. The solid curve is
a DUSTY 100 Myr isochrone from
Chabrier et al. (2000) for masses from
0.1 to 0.03 M .
We took the following steps to identify a set of candidate substellar members of the Pleiades:
keep only objects that fall in the trapezoidal region in Figure 15;
remove objects flagged as nonstellar by the 2MASS pipeline software;
remove objects that appear nonstellar to the eye in the IRAC images;
remove objects that do not fall in or near the locus of field M and L dwarfs in a J - H versus H - Ks diagram;
remove objects that have 3.6 and 4.5 m magnitudes that differ by more than 0.2 mag;
remove objects that fall below the ZAMS in a J versus J - Ks diagram.
As shown in Figure 15, all stars earlier than about mid-M have Ks - [3.6] colors bluer than 0.4. This ensures that
for most of the area of the trapezoidal region, the primary contaminants are distant galaxies. Fortunately, the 2MASS
catalog provides two types of flags for identifying extended objects. For each filter, a 2 flag measures the match between
the objects shape and the instrumental PSF, with values greater than 2.0 generally indicative of a nonstellar object. In
order not to be misguided by an image artifact in one filter, we throw out the most discrepant of the three flags and average
the other two. We discard objects with mean 2 greater than 1.9. The other indicator is the 2MASS extended object flag,
which is the synthesis of several independent tests of the objects shape, surface brightness and color (see Jarrett et al. 2000
for a description of this process). If one simply excludes the objects classified as extended in the 2MASS 6x image by
either of these techniques, the number of candidate VLM and substellar objects lying inside the trapezoidal region
decreases by nearly a half.
We have one additional means to demonstrate that many of the identified objects are probably Pleiades members,
and that is via proper motions. The mean Pleiades proper motion is R.A. = 20 mas yr-1 and decl. = -45 mas yr-1 (Jones
1973). With an epoch difference of only 3.5 yr between the deep 2MASS and IRAC imaging, the expected motion for a
Pleiades member is only 0.07 in right ascension and -0.16 in declination. Given the relatively large pixel size for the
two cameras, and the undersampled nature of the IRAC 3.6 and 4.5 m images, it is not a priori obvious that one would
expect to reliably detect the Pleiades motion. However, both the 2MASS and IRAC astrometric solutions have been very
accurately calibrated. Also, for the present purpose, we only ask whether the data support a conclusion that most of the
identified substellar candidates are true Pleiades members (i.e., as an ensemble), rather than that each star is well enough
separated in a VPD to derive a high membership
probability.
Figure 16 provides a set of plots that we
believe support the conclusion that the majority of
the surviving VLM and substellar candidates are
Pleiades members. The first plot shows the measured
motions between the epoch of the 2MASS and IRAC
observations for all known Pleiades members from
Table 2 that lie in the central square degree region
and have 11 < Ks < 14 (i.e., just brighter than the
substellar candidates). The mean offset of the
Pleiades stellar members from the background
population is well-defined and is quantitatively of the
expected magnitude and sign (+0.07 in right
ascension and -0.16
in declination). The rms
dispersion of the coordinate difference for the field
population in right ascension and declination is 0.076
and 0.062 , supportive of our claim that the
relative astrometry for the two cameras is quite good.
Because we expect that the background population
should have essentially no mean proper motion, the
nonzero mean "motion" of the field population of
about
R.A. = 0.3
is presumably not real.
Instead, the offset is probably due to the uncertainty
in transferring the Spitzer coordinate zero point
between the warm star-tracker and the cryogenic
focal plane. Because it is simply a zero-point offset
applicable to all the objects in the IRAC catalog, it
has no effect on the ability to separate Pleiades
members from the field star population.
Fig. 16 Proper-motion vector point diagrams (VPDs) for various stellar samples in the central 1° field, derived from
combining the IRAC and 2MASS 6x observations. Top left: VPD comparing all objects in the field (small black dots) to
Pleiades members with 11 < Ks < 14 (large blue dots). Top right: Same, except the blue dots are the new candidate VLM
and substellar Pleiades members. Bottom left: Same, except the blue dots are a nearby, low-mass field star sample from
a box just blueward of the trapezoidal region in 15. Bottom right: VPD just showing a second, distant field star sample
as described in the text.
The second panel in Figure 16 shows the proper motion of the candidate Pleiades VLM and substellar objects.
While these objects do not show as clean a distribution as the known members, their mean motion is clearly in the same
direction. After removing 2 deviants, the median offsets for the substellar candidates are 0.04 and -0.11 in right
ascension and declination, respectively. The objects whose motions differ significantly from the Pleiades mean may be
nonmembers or they may be members with poorly determined motions (since a few of the high-probability members in the
first panel also show discrepant motions).
The other two panels in Figure 16 show the proper motions of two possible control samples. The first control
sample was defined as the set of stars that fall up to 0.3 mag below the lower sloping boundary of the trapezoid in Figure
15. These objects should be late-type dwarfs that are either older or more distant than the Pleiades or red galaxies. We used
the 2MASS data to remove extended or blended objects from the sample in the same way as for the Pleiades candidates. If
the objects are nearby field stars, we expect to see large proper motions; if galaxies, the real proper motions would be
small—but relatively large apparent proper motions due to poor centroiding or different centroids at different effective
wavelengths could be present. The second control set was defined to have -0.1 < K - [3.6] < 0.1 and 14.0 < K < 14.5 and to
be stellar based on the 2MASS flags. This control sample should therefore be relatively distant G and K dwarfs primarily.
Both control samples have proper-motion distributions that differ greatly from the Pleiades samples and that make sense
for, respectively, a nearby and a distant field star sample.
Figure 17 shows the Pleiades
members from Table 2 and the 55
candidate VLM and substellar members
that survived all of our culling steps. We
cross-correlated this list with the stars from
Table 2 and with a list of the previously
identified candidate substellar members of
the cluster from other deep imaging
surveys. Fourteen of the surviving objects
correspond to previously identified
Pleiades VLM and substellar candidates.
We provide the new list of candidate
members in Table 5. The columns marked
as (R.A.) and (decl.) are the measured
motions in arcsec over the 3.5 yr epoch
difference between the 2MASS-6x and
IRAC observations. Forty-two of these
objects have Ks > 14.0 and hence inferred
masses less than about 0.1 M ; 31 of them
have Ks > 14.4 and hence have inferred
masses below the hydrogen-burning mass
limit.
(130
kB)
Fig. 17 Same as Fig. 15, except
that the new candidate VLM and
substellar objects from Table 5 are
now indicated as small, red squares.
Table
Candidate
Members
5 New
Pleiades
R.A.
(J2000.0)
(deg)
Decl.
(J2000.0)
(deg)
J
H
Ks
[3.6]
[4.5]
(R.A.)
(decl.)
Previous
ID
SI2M-1...
56.15745
24.42746
14.44
13.79
13.52
13.17
13.10
0.37
-0.01
HHJ 46
SI2M-2...
56.19235
24.38414
14.68
14.10
13.79
13.42
13.36
0.44
-0.19
HHJ 24
SI2M-3...
56.24477
24.27201
17.85
16.83
16.00
15.15
15.15
0.37
0.13
...
SI2M-4...
56.28952
23.97910
15.46
14.83
14.41
14.05
14.05
0.54
-0.16
...
SI2M-5...
56.29098
24.07576
14.80
14.16
13.86
13.43
13.37
0.45
-0.17
...
SI2M-6...
56.30265
23.89584
14.83
14.21
13.88
13.49
13.47
0.37
-0.14
HHJ 14
SI2M-7...
56.32663
23.87112
15.96
15.15
14.79
14.38
14.34
0.18
-0.01
...
SI2M-8...
56.36751
24.52373
16.84
16.05
15.44
...
14.79
0.32
-0.05
...
SI2M-9...
56.39588
23.85472
15.78
15.02
14.65
14.27
14.18
0.37
0.08
...
SI2M-10...
56.40739
23.73057
14.79
14.15
13.81
13.37
13.40
0.33
-0.23
...
SI2M-11...
56.42205
23.90273
15.39
14.73
14.28
13.86
13.85
0.41
-0.10
...
SI2M-12...
56.42644
24.06976
15.27
14.64
14.28
13.89
13.95
0.36
-0.18
...
SI2M-13...
56.43118
23.64760
15.17
14.43
14.14
13.78
13.76
0.36
-0.22
...
SI2M-14...
56.44669
24.51118
17.25
16.37
15.75
...
15.04
0.43
-0.25
...
SI2M-15...
56.45366
23.64644
17.53
16.49
15.57
14.91
14.62
0.21
-0.04
...
SI2M-16...
56.45598
23.95163
14.70
14.07
13.83
13.48
13.36
0.37
-0.08
...
SI2M-17...
56.45634
24.26979
18.11
16.71
16.18
15.38
15.21
0.76
0.10
...
SI2M-18...
56.46099
23.74362
16.39
15.70
15.28
14.64
14.79
0.34
0.01
...
SI2M-19...
56.46113
24.15099
15.81
15.06
14.64
14.08
14.02
0.30
-0.20
BPL 79
SI2M-20...
56.46912
23.86272
15.32
14.64
14.31
13.90
13.78
0.47
-0.18
...
SI2M-21...
56.47910
23.56604
15.57
14.96
14.55
14.18
...
0.44
-0.34
...
SI2M-22...
56.49051
24.05142
14.72
14.12
13.79
13.43
13.44
0.37
-0.11
HHJ 27
SI2M-23...
56.49128
24.41130
16.74
16.09
15.54
14.88
14.82
1.17
0.10
...
SI2M-24...
56.49132
24.14474
14.58
13.99
13.68
13.32
13.32
0.22
-0.06
DH 412
ID
SI2M-25...
56.52133
23.75971
15.56
14.81
14.38
13.98
13.93
0.29
-0.08
...
SI2M-26...
56.52526
23.97200
14.88
14.22
13.93
13.58
13.53
0.26
-0.14
...
SI2M-27...
56.57735
23.98407
14.75
14.11
13.86
13.49
13.39
0.30
-0.08
...
SI2M-28...
56.57843
23.81347
15.83
15.02
14.57
14.06
14.18
0.29
-0.18
...
SI2M-29...
56.58151
23.56235
15.80
15.08
14.69
14.30
...
0.23
-0.14
...
SI2M-30...
56.58557
24.28870
17.05
16.31
15.76
15.09
14.87
-0.10
-0.12
...
SI2M-31...
56.59283
23.87408
15.59
14.94
14.46
14.10
14.01
0.45
-0.11
...
SI2M-32...
56.60060
24.50354
14.75
14.14
13.84
13.48
13.42
0.38
-0.22
BPL
101
SI2M-33...
56.60880
24.08598
15.19
14.52
14.15
13.72
13.74
0.40
-0.07
...
SI2M-34...
56.63392
24.38740
17.25
16.34
15.77
15.13
15.11
0.27
-0.23
...
SI2M-35...
56.64737
23.95206
15.37
14.77
14.45
14.06
13.97
0.29
-0.22
...
SI2M-36...
56.67914
24.41405
15.55
14.85
14.42
13.98
14.00
0.26
-0.17
BPL
108
SI2M-37...
56.70850
24.00659
15.69
14.98
14.58
14.05
14.00
0.34
-0.04
...
SI2M-38...
56.75776
24.22451
14.97
14.41
14.10
13.73
13.62
0.29
-0.06
BPL
122
SI2M-39...
56.77373
24.66767
15.47
14.87
14.51
...
14.01
0.36
-0.07
...
SI2M-40...
56.79400
23.90606
15.99
15.28
15.02
14.57
14.54
0.19
0.01
...
SI2M-41...
56.79446
23.97119
14.95
14.38
14.02
13.66
13.66
0.32
-0.20
...
SI2M-42...
56.79918
24.22539
14.86
14.23
13.88
13.49
13.38
0.27
-0.08
BPL
130
SI2M-43...
56.80051
24.47547
16.19
15.53
15.05
14.57
14.65
0.41
-0.41
BPL
132
SI2M-44...
56.82203
24.20922
17.54
17.00
15.94
15.23
14.62
-0.05
-0.10
...
SI2M-45...
56.96009
23.91330
16.41
15.71
15.20
14.51
14.53
0.34
-0.16
...
SI2M-46...
56.96365
23.73669
17.52
16.67
16.02
15.22
15.11
0.25
-0.04
...
SI2M-47...
57.00899
24.42107
16.75
16.05
15.42
14.85
14.76
0.23
-0.08
...
SI2M-48...
57.01952
23.65838
15.28
14.66
14.27
13.78
...
0.26
-0.15
...
SI2M-49...
57.07928
24.42024
16.02
15.27
14.95
14.49
14.51
0.36
-0.11
BPL
172
SI2M-50...
57.09851
24.37646
14.92
14.34
13.94
13.58
13.55
0.30
-0.01
BPL
177
SI2M-51...
57.12811
23.70665
16.66
15.85
15.19
14.57
...
0.32
-0.12
...
SI2M-52...
57.13138
24.57707
16.78
15.88
15.38
14.70
14.65
0.18
-0.02
...
SI2M-53...
57.14174
24.08293
15.38
14.67
14.47
14.10
14.00
0.44
-0.16
...
SI2M-54...
57.23196
24.36115
15.04
14.43
14.10
13.69
13.66
0.36
-0.25
HHJ 8
SI2M-55...
57.28922
23.94612
17.70
16.77
16.09
15.14
15.02
0.54
0.27
...
Our candidate list could be contaminated by foreground late-type dwarfs that happen to lie in the line of sight to
the Pleiades. How many such objects should we expect? In order to pass our culling steps, such stars would have to be
mid- to late-M dwarfs, or early to mid-L dwarfs. We use the known M dwarfs within 8 pc to estimate how many field M
dwarfs should lie in a 1 deg2 region and at distance between 70 and 100 pc (so they would be coincident in a CMD with the
100 Myr Pleiades members). The result is 3 such field M dwarf contaminants. Cruz et al. (2007) estimate that the
volume density of L dwarfs is comparable to that for late-M dwarfs, and therefore a very conservative estimate is that there
might also be 3 field L dwarfs contaminating our sample. We regard this (6 contaminating field dwarfs) as an upper limit
because our various selection criteria would exclude early-M dwarfs and late-L dwarfs. Bihain et al. (2006) made an
estimate of the number of contaminating field dwarfs in their Pleiades survey of 1.8 deg2; for the spectral type range of our
objects, their algorithm would have predicted just one or two contaminating field dwarfs for our survey.
How many substellar Pleiades members should there be in the region we have surveyed? That is, of course, part
of the question we are trying to answer. However, previous studies have estimated that the Pleiades stellar mass function
for M < 0.5 M can be approximated as a power law with an exponent of -1 (dN/dM M-1). Using the known Pleiades
members from Table 2 that lie within the region of the IRAC survey and that have masses of 0.2 < M/M < 0.5 (as
estimated from the Baraffe et al. (1998) 100 Myr isochrone) to normalize the relation, the M-1 mass function predicts about
48 members in our search region and with 14 < K < 16.2 (corresponding to 0.1 < M/M < 0.035). Other studies have
suggested that the mass function in the Pleiades becomes shallower below 0.1 M , dN/dM M-0.6. Using the same
normalization as above, this functional form for the Pleiades mass function for M < 0.1 M yields a prediction of 20
VLM and substellar members in our survey. The number of candidates we have found falls between these two estimates.
Better proper motions and low-resolution spectroscopy will almost certainly eliminate some of these candidates as
nonmembers.
6. MID-IR OBSERVATIONS OF DUST AND POLYCYCLIC AROMATIC HYDROCARBONS IN THE
PLEIADES
Since the earliest days of astrophotography, it has been clear that the Pleiades stars are in relatively close
proximity to interstellar matter whose optical manifestation is the spider-web–like network of filaments seen particularly
strongly toward several of the B stars in the cluster. High-resolution spectra of the brightest Pleiades stars as well as CO
maps toward the cluster show that there is gas as well as dust present and that the (primary) interstellar cloud has a
significant radial velocity offset relative to the Pleiades (White 2003; Federman & Willson 1984). The gas and dust,
therefore, are not a remnant from the formation of the cluster but are simply evidence of a transitory event as this small
cloud passes by the cluster in our line of sight (see also Breger 1986). There are at least two claimed morphological
signatures of a direct interaction of the Pleiades with the cloud. White & Bally (1993) provided evidence that the IRAS 60
and 100 m image of the vicinity of the Pleiades showed a dark channel immediately to the east of the Pleiades, which
they interpreted as the "wake" of the Pleiades as it plowed through the cloud from the east. Herbig & Simon (2001)
provided a detailed analysis of the optically brightest nebular feature in the Pleiades—IC 349 (Barnard's Merope
nebula)—and concluded that the shape and structure of that nebula could best be understood if the cloud was running into
the Pleiades from the southeast. Herbig & Simon (2001) concluded that the IC 349 cloudlet, and by extension the rest of
the gas and dust enveloping the Pleiades, are relatively distant outliers of the Taurus molecular clouds (see also Eggen
1950 for a much earlier discussion ascribing the Merope nebulae as outliers of the Taurus clouds). White (2003) has more
recently proposed a hybrid model, where there are two separate interstellar cloud complexes with very different space
motions, both of which are colliding simultaneously with the Pleiades and with each other.
Breger (1986) provided polarization measurements for a sample of member and background stars toward the
Pleiades and argued that the variation in polarization signatures across the face of the cluster was evidence that some of the
gas and dust was within the cluster. In particular, Figure 6 of that paper showed a fairly distinct interface region, with little
residual polarization to the NE portion of the cluster and an L-shaped boundary running EW along the southern edge of the
cluster and then north-south along the western edge of the cluster. Stars to the south and west of that boundary show
relatively large polarizations and consistent angles (see also our Fig. 5, where we provide a few polarization vectors from
Breger 1986 to illustrate the location of the interface region and the fact that the position angle of the polarization
correlates well with the location in the interface).
There is a general correspondence between the polarization map and what is seen with IRAC, in the sense that
the B stars in the NE portion of the cluster (Atlas and Alcyone) have little nebular emission in their vicinity, whereas those
in the western part of the cluster (Maia, Electra, and Asterope) have prominent, filamentary dust emission in their vicinity.
The L-shaped boundary is in fact visible in Figure 4 as enhanced nebular emission running between and below a line
roughly joining Merope and Electra and then making a right angle and running roughly parallel to a line running from
Electra to Maia to HII 1234 (see Fig. 5).
6.1. Pleiades Dust-Star Encounters Imaged with IRAC
The Pleiades dust filaments are most strongly evident in IRAC's 8 m channel, as evidenced by the distinct red
color of the nebular features in Figure 4. The dominance at 8 m is an expected feature of reflection nebulae, as
exemplified by NGC 7023 (Werner et al. 2004), where most of the mid-infrared emission arises from polycyclic aromatic
hydrocarbons (PAHs) whose strongest bands in the 3–10 m region fall at 7.7 and 8.6 m. One might expect that if
portions of the passing cloud were particularly near to one of the Pleiades members, it might be possible to identify such
interactions by searching for stars with 8.0 m excesses or for stars with extended emission at 8 m. Figure 18 provides
two such plots. Four stars stand out
as having significant extended 8 m
emission, with two of those stars also
having an 8 m excess based on
their [3.6] - [8.0] color. All of these
stars, plus IC 349, are located
approximately along the interface
region identified by Breger (1986).
Fig. 18 Two plots intended to
(67 isolate Pleiades members with
kB) excess and/or extended 8 m
emission. The plot with [3.6] [8.0] m colors shows data
from Table 3 (and hence is for
aperture sizes of 3 pixel and 2
pixel radius, respectively). The
increased vertical spread in the
plots at faint magnitudes is
simply due to decreasing
signal-to-noise at 8 m. The
numbers labeling stars with
excesses
are
the
HII
identification numbers for
those stars.
We have subtracted a PSF from the 8 m images for the stars with extended emission, and those PSF-subtracted
images are provided in Figure 19. The image for HII 1234 has the appearance of a bow shock. The shape is reminiscent of
predictions for what one should expect from a collision between a large cloud or a sheet of gas and an A star as described
in Artymowicz & Clampin (1997). The Artymowicz & Clampin model posits that A stars encountering a cloud will carve
a paraboloidal shaped cavity in the cloud via radiation pressure. The exact size and shape of the cavity depend on the
relative velocity of the encounter, the star's mass and luminosity and properties of the ISM grains. For typical parameters,
the predicted characteristic size of the cavity is of order 1000 AU, quite comparable to the size of the structures around HII
652 and HII 1234. The observed appearance of the cavity depends on the view angle to the observer. However, in any
case, the direction from which the gas is moving relative to the star can be inferred from the location of the star relative to
the curved rim of the cavity; the "wind" originates approximately from the direction connecting the star and the apex of the
rim. For HII 1234, this indicates the cloud that it is encountering has a motion relative to HII 1234 from the SSE, in accord
with a Taurus origin and not in accord for where a cloud is impacting the Pleiades from the west as posited in White
(2003). The nebular emission for HII 652 is less strongly bow-shaped, but the peak of the excess emission is displaced
roughly southward from the star, consistent with the Taurus model and inconsistent with gas flowing from the west.
Fig. 19 Postage stamp images
extracted from individual, 8 m
BCDs for the stars with extended 8
m emission, from which we have
subtracted an empirical PSF.
Clockwise from the upper left, the
stars shown are HII 1234, HII 859,
Merope, and HII 652. The fivepointed
star
indicates
the
astrometric position of the star
(often superposed on a few black
pixels where the 8 m image was
saturated. The circle in the Merope
image is centered on the location of
IC 349 and has a diameter of about
25 (the size of IC 349 in the
optical is of order 10 × 10 ).
Despite being the brightest part of the Pleiades nebulae in the optical, IC 349 appears to be undetected in the 8
m image. This is not because the 8 m image is insensitive to the nebular emission—there is generally good agreement
between the structures seen in the optical and at 8 m, and most of the filaments present in optical images of the Pleiades
are also visible on the 8 m image (see Figs. 4 and 19) and even the PSF-subtracted image of Merope shows well-defined
nebular filaments. The lack of enhanced 8 m emission from the region of IC 349 is probably because all of the small
particles have been scoured away from this cloudlet, consistent with Herbig's model to explain the HST surface
photometry and colors. There is no PAH emission from IC 349 because there are none of the small molecules that are the
postulated source of the PAH emission.
IC 349 is very bright in the optical, and undetected to a good sensitivity limit at 8 m; it must be detectable via
imaging at some wavelength between 5000 Å and 8 m. We checked our 3.6 m data for this purpose. In the standard
BCD mosaic image, we were unable to discern an excess at the location of IC 349 either simply by displaying the image
with various stretches or by doing cuts through the image. We performed a PSF subtraction of Merope from the image in
order to attempt to improve our ability to detect faint, extended emission 30 from Merope—unfortunately, bright stars
have ghost images in IRAC channel 1, and in this case the ghost image falls almost exactly at the location of IC 349. IC
349 is also not detected in visual inspection of our 2MASS 6x images.
6.2. Circumstellar Disks and IRAC
As part of the Spitzer FEPS (Formation and Evolution of Planetary Systems) Legacy program, using pointed
MIPS photometry, Stauffer et al. (2005) identified three G dwarfs in the Pleiades as having 24 m excesses probably
indicative of circumstellar dust disks. Gorlova et al. (2006) reported results of a MIPS GTO survey of the Pleiades and
identified nine cluster members that appear to have 24 m excesses due to circumstellar disks. However, it is possible that
in a few cases these apparent excesses could be due instead to a knot of the passing interstellar dust impacting the cluster
member or that the 24 m excess could be flux from a background galaxy projected onto the line of sight to the Pleiades
member. Careful analysis of the IRAC images of these cluster members may help confirm that the MIPS excesses are
evidence for debris disks rather than the other possible explanations.
Six of the Pleiades members with probable 24 m excesses are included in the region mapped with IRAC.
However, only four of them have data at 8 m—the other two fall near the edge of the mapped region and only have data
at 3.6 and 5.8 m. None of the six stars appear to have significant local nebular dust from visual inspection of the IRAC
mosaic images. Also, none of them appear problematic in Figure 18. For a slightly more quantitative analysis of possible
nebular contamination, we also constructed aperture growth curves for the six stars and compared them to other Pleiades
members. All but one of the six show aperture growth curves that are normal and consistent with the expected IRAC PSF.
The one exception is HII 489, which has a slight excess at large aperture sizes, as is illustrated in Figure 20. Because HII
489 only has a small 24 m excess, it is possible that the 24 m excess is due to a local knot of the interstellar cloud
material and is not due to a debris disk. For the other five 24 m excess stars we find no such problem, and we conclude
that their 24 m excesses are indeed best explained as due to debris disks.
Fig. 20 Aperture growth curves from the 8
m mosaic for stars with 24 m excesses from
Gorlova et al. (2006) and for a set of control
objects (dashed curves). All of the objects have
been scaled to common zero-point magnitudes
for 9 pixel apertures, with the 24 m excess
stars offset from the control objects by 0.1 mag.
The three Gorlova et al. (2006) stars with no
excess at 8 m are HII 996, HII 1284, and HII
2195. The Gorlova et al. (2006) star with a
slight excess at 8 m is HII 489.
(52 kB)
7. SUMMARY AND CONCLUSIONS
We have collated the primary membership catalogs for the Pleiades to produce the first catalog of the cluster
extending from its highest mass members to the substellar limit. At the bright end, we expect this catalog to be essentially
complete and with few or no nonmember contaminants. At the faint end, the data establishing membership are much
sparser, and we expect a significant number of objects will be nonmembers. We hope that the creation of this catalog will
spur efforts to obtain accurate radial velocities and proper motions for the faint candidate members in order to eventually
provide a well-vetted membership catalog for the stellar members of the Pleiades. Toward that end, it would be useful to
update the current catalog with other data—such as radial velocities, lithium equivalent widths, X-ray fluxes, H
equivalent widths, etc.—which could be used to help accurately establish membership for the low-mass cluster candidates.
It is also possible to make more use of "negative information" present in the proper-motion catalogs. That is, if a member
from one catalog is not included in another study but does fall within its areal and luminosity coverage, that suggests that it
likely failed the membership criteria of the second study. For a few individual stars, we have done this type of comparison,
but a systematic analysis of the proper-motion catalogs should be conducted. We intend to undertake these tasks and plan
to establish a Web site where these data would be hosted.
We have used the new Pleiades member catalog to define the single-star locus at 100 Myr for BVICKs and the
four IRAC bands. These curves can be used as empirical calibration curves when attempting to identify members of less
well-studied, more distant clusters of similar age. We compared the Pleiades photometry to theoretical isochrones from
Siess et al. (2000) and Baraffe et al. (1998). The Siess et al. (2000) isochrones are not, in detail, a good fit to the Pleiades
photometry, particularly for low-mass stars. The Baraffe et al. (1998) 100 Myr isochrone does fit the Pleiades photometry
very well in the I versus I - K plane.
We have identified 31 new substellar candidate members of the Pleiades using our combined seven-band infrared
photometry and have shown that the majority of these objects appear to share the Pleiades proper motion. We believe that
most of the objects that may be contaminating our list of candidate brown dwarfs are likely to be unresolved galaxies, and
therefore low-resolution spectroscopy should be able to provide a good criterion for culling our list of nonmembers.
The IRAC images, particularly the 8 m mosaic, provide vivid evidence of the strong interaction of the Pleiades
stars and the interstellar cloud that is passing through the Pleiades. Our data are supportive of the model proposed by
Herbig & Simon (2001) whereby the passing cloud is part of the Taurus cloud complex and hence is encountering the
Pleiades from the SSE direction. White & Bally (1993) had proposed a model whereby the cloud was encountering the
Pleiades from the west and used this to explain features in the IRAS 60 and 100 m images of the region as the wake of
the Pleiades moving through the cloud. Our data appear to not be supportive of that hypothesis and therefore leave the
apparent structure in the IRAS maps as unexplained.
Most of the support for this work was provided by the Jet Propulsion Laboratory, California Institute of
Technology, under NASA contract 1407. This research has made use of NASA's Astrophysics Data System (ADS)
Abstract Service, and of the SIMBAD database, operated at CDS, Strasbourg, France. This research has made use of data
products from the Two Micron All-Sky Survey (2MASS), which is a joint project of the University of Massachusetts and
the Infrared Processing and Analysis Center, funded by the National Aeronautics and Space Administration and the
National Science Foundation. These data were served by the NASA/IPAC Infrared Science Archive, which is operated by
the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space
Administration. The research described in this paper was partially carried out at the Jet Propulsion Laboratory, California
Institute of Technology, under contract with the National Aeronautics and Space Administration.
This research made use of the SIMBAD database operated at CDS, Strasbourg, France, and also of the NED and
NStED databases operated at IPAC, Pasadena, CA. A large amount of data for the Pleiades (and other open clusters) can
also be found at the open cluster database WEBDA (http://www.univie.ac.at/webda/), operated in Vienna by Ernst
Paunzen.
APPENDIX
A1. MEMBERSHIP CATALOGS
Membership lists of the Pleiades date back to antiquity if one includes historical and literary references to the
Seven Sisters (Alcyone, Maia, Merope, Electra, Taygeta, Asterope, and Celeno) and their parents (Atlas and Pleione). The
first paper discussing relative proper motions of a large sample of stars in the Pleiades (based on visual observations) was
published by Pritchard (1884). The best of the early proper-motion surveys of the Pleiades derived from photographic
plate astrometry was that by Trumpler (1921), based on plates obtained at Yerkes and Lick observatories. The candidate
members from that survey were presented in two tables, with the first being devoted to candidate members within about 1°
from the cluster center (operationally, within 1° from Alcyone) and the second table being devoted to candidates further
than 1° from the cluster center. Most of the latter stars were denoted by Trumpler by an S or R, followed by an
identification number. We use Tr to designate the Trumpler stars (hence Trnnn for a star from the first table and the small
number of stars in the second table without an "S" or an "R," and TrSnnn or TrRnnn for the other stars). For the central
region, Trumpler's catalog extends to V 13, while the outer region catalog includes stars only to about V 9.
The most heavily referenced proper-motion catalog of the Pleiades is that provided by Hertzsprung (1947). That
paper makes reference to two separate catalogs: a photometric catalog of the Pleiades published by Hertzsprung (1923),
whose members are commonly referred to by HI numbers, and the new proper-motion catalog from the 1947 paper,
commonly referenced as the HII catalog. While both HI and HII numbers have been used in subsequent observational
papers, it is the HII identification numbers that predominate. That catalog—derived from Carte du Ciel blue-sensitive
plates from 14 observatories—includes stars in the central 2 × 2 region of the cluster and has a faint limit of about V =
15.5. Johnson system BVI photometry is provided for most of the proposed Hertzsprung members in Johnson & Mitchell
(1958) and Iriarte (1967). Additional Johnson B and V photometry plus Kron I photometry for a fairly large number of the
Hertzsprung members can be found in Stauffer (1980, 1982, 1984). Other Johnson BV photometry for a scattering of stars
can be found in Jones (1973), Robinson & Kraft (1974), and Messina (2001). Spectroscopic confirmation, primarily via
radial velocities, that these are indeed Pleiades members has been provided in Soderblom et al. (1993), Queloz et al.
(1998), and Mermilliod et al. (1997).
Two other proper-motion surveys provide relatively bright candidate members relatively far from the cluster
center: Artyukhina & Kalinina (1970) and van Leeuwen 1986. Stars from the Artyukhina catalog are designated as "AK"
followed by the region from which the star was identified followed by an identification number. The new members
provided in the van Leeuwen paper were taken from an otherwise unpublished proper-motion study by Pels, where the
first 118 stars were considered probable members and the remaining 75 stars were considered possible members. Van
Leeuwen categorized a number of the Pels stars as nonmembers based on the Walraven photometry they obtained, and we
adopt those findings. Radial velocities for stars in these two catalogs have been obtained by Rosvick et al. (1992),
Mermilliod et al. (1997), and Queloz et al. (1998), and those authors identified a list of the candidate members that they
considered confirmed by the high-resolution spectroscopy. For these outlying candidate members, to be included in Table
2 we require that the star be a radial velocity member from one of the above three surveys, or be indicated as having "no
dip" in the Coravel cross-correlation (indicating rapid rotation, which at least for the later type stars is suggestive of
membership). Geneva photometry of the Artyukhina stars considered as likely members was provided by Mermilliod et al.
(1997). The magnitude limit of these surveys was not well-defined, but most of the Artyukhina and Pels stars are brighter
than V = 13.
Jones (1973) provided proper-motion membership probabilities for a large sample of proposed Pleiades members,
and for a set of faint, red stars toward the Pleiades. A few star identification names from the sources considered by Jones
appear in Table 2, including MT (McCarthy & Treanor 1964), VM (van Maanen 1946), and ALR (Ahmed et al. 1965;
Jones 1973).
The chronologically next significant source of new Pleiades candidate members was the flare star survey of the
Pleiades conducted at several observatories in the 1960s, and summarized in Haro et al. (1982, hereafter HCG). The logic
behind these surveys was that even at 100 Myr, late-type dwarfs have relatively frequent and relatively high-luminosity
flares (as demonstrated by Johnson & Mitchell 1958 having detected two flares during their photometric observations of
the Pleiades), and therefore wide area, rapid cadence imaging of the Pleiades at blue wavelengths should be capable of
identifying low-mass cluster members. However, such surveys also will detect relatively young field dwarfs, and therefore
it is best to combine the flare star surveys with proper motions. Dedicated proper-motion surveys of the HCG flare stars
were conducted by Jones (1981) and Stauffer et al. (1991), with the latter also providing photographic VI photometry
(Kron system). Photoelectric photometry for some of the HCG stars have been reported in Stauffer (1982, 1984), Stauffer
& Hartmann (1987), and Prosser et al. (1991). High-resolution spectroscopy of many of the HCG stars is reported in
Stauffer (1984), Stauffer & Hartmann (1987), and Terndrup et al. (2000). Because a number of the papers providing
additional observational data for the flare stars were obtained prior to 1982, we also include in Table 2 the original flare
star names that were derived from the observatory where the initial flare was detected. Those names are of the form of an
initial letter indicating the observatory—A (Asiago), B (Byurakan), K (Konkoly), T (Tonantzintla)—followed by an
identification number.
Stauffer et al. (1991) conducted two proper-motion surveys of the Pleiades over an approximately 4 × 4 region
of the cluster based on plates obtained with the Lick 20 astrographic telescope. The first survey was essentially unbiased,
except for the requirement that the stars fall approximately in the region of the V versus V - I color-magnitude diagram
where Pleiades members should lie. Candidate members from this survey are designated by SK numbers. The second
survey was a proper-motion survey of the HCG stars. Photographic VI photometry of all the stars was provided as well as
proper-motion membership probabilities. Photoelectric photometry for some of the candidate members was obtained as
detailed above in the section on the HCG catalog stars. The faint limit of these surveys is about V = 18.
Hambly et al. (1991) provided a significantly deeper, somewhat wider area proper-motion survey, with the
faintest members having V 20 and the total area covered being of order 25 deg2. The survey utilized red sensitive plates
from the Palomar and UK Schmidt telescopes. Due to incomplete coverage at one epoch, there is a vertical swath slightly
east of the cluster center where no membership information is available. Stars from this survey are designated by their HHJ
numbers. Hambly et al. (1993) provide RI photographic photometry on a natural system for all of their candidate members,
plus photoelectric Cousins RI photometry for a small number of stars and JHK photometry for a larger sample. Some
spectroscopy to confirm membership has been reported in Stauffer et al. (1994, 1995, 1999), Oppenheimer et al. (1997),
and Steele et al. (1995), although for most of the HHJ stars there is no spectroscopic membership confirmation.
Pinfield et al. (2000) provide the deepest wide-field proper-motion survey of the Pleiades. That survey combines
CCD imaging of 6 deg2 of the Pleiades obtained with the Burrell Schmidt telescope (as five separate, nonoverlapping
fields near but outside the cluster center) with deep photographic plates that provide the first epoch positions. Candidate
members are designated by BPL numbers (for Burrell Pleiades), with the faintest stars having I 19.5, corresponding to
V > 23. Only the stars brighter than about I = 17 have sufficiently accurate proper motions to use to identify Pleiades
members. Fainter than I = 17, the primary selection criteria are that the star fall in an appropriate place in both an I versus I
- Z and an I versus I - K CMD.
Adams et al. (2001) combined the 2MASS and digitized POSS databases to produce a very wide area propermotion survey of the Pleiades. By design, that survey was very inclusive—covering the entire physical area of the cluster
and extending to the hydrogen-burning mass limit. However, it was also very "contaminated," with many suspected
nonmembers. The catalog of possible members was not published. We have therefore not included stars from this study in
Table 2; we have used the proper-motion data from Adams et al. (2001) to help decide cases where a given star has
ambiguous membership data from the other surveys.
Deacon & Hambly (2004) provided another deep and very wide area proper-motion survey of the Pleiades. The
survey covers a circular area of approximately 5° radius to R 20, or V 22. Candidate members are designated by
"DH." Deacon & Hambly (2004) also provide membership probabilities based on proper motions for many candidate
cluster members from previous surveys. For stars where Deacon & Hambly (2004) derive P < 0.1 and where we have no
other proper-motion information or where another proper-motion survey also finds low membership probability, we
exclude the star from our catalog. For cases where two of our proper-motion catalogs differ significantly in their
membership assessment, with one survey indicating the star is a probable member, we retain the star in the catalog as the
conservative choice. Examples of the latter where Deacon & Hambly (2004) derive P < 0.1 include HII 1553, HII 2147,
HII 2278, and HII 2665—all of which we retain in our catalog because other surveys indicate these are high-probability
Pleiades members.
A2. PHOTOMETRY
Photometry for stars in open cluster catalogs can be used to help confirm cluster membership and to help
constrain physical properties of those stars or of the cluster. For a variety of reasons, photometry of stars in the Pleiades
has been obtained in a panoply of different photometric systems. For our own goals, which are to use the photometry to
help verify membership and to define the Pleiades single-star locus in color-magnitude diagrams, we have attempted to
convert photometry in several of these systems to a common system (Johnson BV and Cousins I). We detail below the
sources of the photometry and the conversions we have employed.
Photoelectric photometry of Pleiades members dates back to at least 1921 Cummings (1921). However, as far as
we are aware the first "modern" photoelectric photometry for the Pleiades, using a potassium hydride photoelectric cell, is
that of Calder & Shapley (1937). Eggen (1950) provided photoelectric photometry using a 1P21 phototube (but calibrated
to a no-longer-used photographic system) for most of the known Pleiades members within 1° of the cluster center and with
magnitudes <11. The first phototube photometry of Pleiades stars calibrated more-or-less to the modern UBV system was
provided by Johnson & Morgan (1951). An update of that paper, and the oldest photometry included here was reported in
Johnson & Mitchell (1958), which provided UBV Johnson system photometry for a large sample of HII and Trumpler
candidate Pleiades members. Iriarte (1967) later reported Johnson system V - I colors for most of these stars. We have
converted Iriarte's V - I photometry to estimated Cousins V - I colors using a formula from Bessell (1979):
BVRI photometry for most of the Hertzsprung members fainter than V = 10 has been published by Stauffer (1980,
1982, 1984) and Stauffer & Hartmann (1987). The BV photometry is Johnson system, whereas the RI photometry is on the
Kron system. The Kron V - I colors were converted to Cousins V - I using a transformation provided by Bessell & Weis
(1987):
Other Kron system V - I colors have been published for Pleiades candidates in Stauffer et al. (1991, photographic
photometry) and in Prosser et al. (1991). These Kron-system colors have also been converted to Cousins V - I using the
above formula.
Johnson/Cousins UBVR photometry for a set of low-mass Pleiades members was provided by Landolt (1979).
We only use the BV magnitudes from that study. Additional Johnson system UBV photometry for small numbers of stars is
provided in Robinson & Kraft (1974), Messina (2001), and Jones (1973).
Van Leeuwen et al. (1987) provided Walraven VBLUW photometry for nearly all of the Hertzsprung members
brighter than V 13.5 and for the Pels candidate members. Van Leeuwen provided an estimated Johnson V derived from
the Walraven V in his tables. We have transformed the Walraven V - B color into an estimate of Johnson B - V using a
formula
from
Rosvick
et
al.
(1992):
Hambly et al. (1993) provided photographic VRI photometry for all of the HHJ candidate members and VRI
Cousins photoelectric photometry for a small fraction of those stars. We took all of the HHJ stars with photographic
photometry for which we also have photoelectric VI photometry on the Cousins system, and plotted V(Cousins) versus
V(HHJ) and I(Cousins) versus I(HHJ). While there is some evidence for slight systematic departures of the HHJ
photographic photometry from the Cousins system, those departures are relatively small and we have chosen simply to
retain the HHJ values and treat them as Cousins system.
Pinfield et al. (2000) reported their I magnitudes in an instrumental system that they designated as Ikp. We
identified all BPL candidate members for which we had photoelectric Cousins I estimates, and plotted Ikp versus IC. Figure
21 shows this correlation, and the piecewise linear fit we have made to convert from Ikp to IC. Our catalog lists these
converted IC measures for the BPL stars for which we have no other photoelectric I estimates.
Fig. 21 Calibration derived relating Ikp
from Pinfield et al. (2000) and IC. The
dots represent stars for which we have
both Ikp and IC measurements (small dots:
photographic IC; large dots: photoelectric
IC), and the solid line indicates the
piecewise linear fit we use to convert the
Ikp values to IC for stars for which we only
have Ikp.
Deacon & Hambly (2004) derived RI photometry from the scans of their plates and calibrated that photometry by
reference to published photometry from the literature. When we plotted their the difference between their I-band
photometry and literature values (where available), we discovered a significant dependence on right ascension.
Unfortunately, because the DH survey extended over larger spatial scales than the calibrating photometry, we could not
derive a correction that we could apply to all the DH stars. We therefore developed the following indirect scheme. We
used the stars for which we have estimated IC magnitudes (from photoelectric photometry) to define the relation between J
and (IC - J) for Pleiades members. For each DH star, we combined that relation and the 2MASS J magnitude to yield a
predicted IC. Figure 22 shows a plot of the difference of this predicted IC and I(DH) with right ascension. The solid line
shows the relation we adopt. Figure 23 shows the relation between the corrected I(DH) values and Table 2 IC measures
from photoelectric sources. There is still a significant amount of scatter, but the corrected I(DH) photometry appears to be
accurately calibrated to the Cousins system.
Fig. 22 Difference between the predicted IC
and Deacon & Hambly (2004) I magnitude as
a function of right ascension for the DH stars.
No obvious dependence is present vs.
declination.
Fig. 23 Comparison of the recalibrated DH I photometry with estimates of IC for stars in Table 2 with photoelectric
data
In a very few cases (specifically, just five stars), we provide an estimate of IC based on data from a wide-area
CCD survey of Taurus obtained with the Quest-2 camera on the Palomar 48 inch Samuel Oschin telescope (Slesnick et al.
2006). That survey calibrated their photometry to the Sloan i system, and we have converted the Sloan i magnitudes to IC.
We intend to make more complete use of the Quest-2 data in a subsequent paper.
When we have multiple sources of photometry for a given star, we consider how to combine them. In most cases,
if we have photoelectric data, that is given preference. However, if we have photographic V and I, and only a photoelectric
measurement for I, we do not replace the photographic I with the photoelectric value because these stars are variable and
the photographic measurements are at least in some cases from nearly simultaneous exposures. Where we have multiple
sources for photoelectric photometry, and no strong reason to favor one measurement or set of measurements over another,
we have averaged the photometry for a given star. In most cases where we have multiple photometry the individual
measurements agree reasonably well but with the caveat that the Pleiades low-mass stars are in many cases heavily spotted
and "active" chromospherically and hence are photometrically variable. In a few cases, even given the expectation that
spots and other phenomena may affect the photometry, there seems to be more discrepancy between reported V
magnitudes than we expect. We note two such cases here. We suspect these results indicate that at least some of the
Pleiades low-mass stars have long-term photometric variability larger than their short period (rotational) modulation.
HII 882 has at least four presumably accurate V magnitude measurements reported in the literature. Those
measures are V = 12.66 Johnson & Mitchell (1958); V = 12.95 Stauffer (1982); V = 12.898 van Leeuwen et al. (1986); and
V = 12.62 Messina (2001).
HII 345 has at least three presumably accurate V magnitude measurements. Those measurements are V = 11.65
Landolt (1979); V = 11.73 van Leeuwen et al. (1986); V = 11.43 Messina (2001).
At the bottom of Table 2, we provide a key to the source(s) of the optical photometry provided in the table.
REFERENCES
Adams, J., Stauffer, J., Monet, D., Skrutskie, M., & Beichman, C. 2001, AJ, 121, 2053 First citation in article |
IOPscience | ADS
Ahmed, F., Lawrence, L., & Reddish, V. 1965, Publ. R. Obs. Edinburgh, 3, 187 First citation in article | ADS
Allen, L., et al. 2004, ApJS, 154, 363 First citation in article | IOPscience | ADS
Artymowicz, P., & Clampin, M. 1997, ApJ, 490, 863 First citation in article | IOPscience | ADS
Artyukhina, N. 1969, Soviet Astron., 12, 987 First citation in article | ADS
Artyukhina, N., & Kalinina, E. 1970, Trudy Sternberg Astron. Inst., 40, 3 First citation in article | ADS
Baraffe, I., Chabrier, G., Allard, F., & Hauschildt, P. 1998, A&A, 337, 403 First citation in article | ADS
Bessell, M. 1979, PASP, 91, 589 First citation in article | CrossRef | ADS
Bessell, M., & Weis, E. 1987, PASP, 99, 642 First citation in article | CrossRef | ADS
Bihain, G., et al. 2006, A&A, 458, 805 First citation in article | CrossRef | ADS
Breger, M. 1986, ApJ, 309, 311 First citation in article | CrossRef | ADS
Calder, W., & Shapley, H. 1937, Ann. Ast. Obs. Harvard College, 105, 453 First citation in article | ADS
Chabrier, G., Baraffe, I., Allard, F., & Hauschildt, P. 2000, ApJ, 542, 464 First citation in article | IOPscience | ADS
Cruz, K., et al. 2007, AJ, 133, 439 First citation in article | IOPscience | ADS
Cummings, E. 1921, PASP, 33, 214 First citation in article | CrossRef | ADS
Deacon, N., & Hambly, N. 2004, A&A, 416, 125 First citation in article | CrossRef | ADS
Eggen, O. 1950, ApJ, 111, 81 First citation in article | CrossRef | ADS
Federman, S., & Willson, R. 1984, ApJ, 283, 626 First citation in article | CrossRef | ADS
Festin, L. 1998, A&A, 333, 497 First citation in article | ADS
Gorlova, N., et al. 2006, ApJ, 649, 1028 First citation in article | IOPscience | ADS
Hambly, N., Hawkins, M. R. S., & Jameson, R. 1991, MNRAS, 253, 1 First citation in article | ADS
———. 1993, A&AS, 100, 607 First citation in article | ADS
Haro, G., Chavira, E., & Gonzalez, G. 1982, Bol. Inst. Tonantzintla, 3, 1 First citation in article | ADS
Herbig, G., & Simon, T. 2001, AJ, 121, 3138 First citation in article | IOPscience | ADS
Hertzsprung, E. 1923, Mem. Danish Acad. 4, 4 First citation in article
———. 1947, Ann. Leiden Obs., 191, 1 First citation in article | ADS
Iriarte, B. 1967, Bol. Obs. Tonantzintla Tacubaya, 4, 79 First citation in article | ADS
Jameson, R., & Skillen, I. 1989, MNRAS, 239, 247 First citation in article | ADS
Jarrett, T., Chester, T., Cutri, R., Schneider, S., Skrutskie, M., & Huchra, J. 2000, AJ, 119, 2498 First citation in article |
IOPscience | ADS
Jeffries, R. D., & Oliveira, J. 2005, MNRAS, 358, 13 First citation in article | CrossRef | ADS
Johnson, H. L., & Mitchell, R. I. 1958, ApJ, 128, 31 (JM) First citation in article | CrossRef | ADS
Johnson, H. L., & Morgan, W. W. 1951, ApJ, 114, 522 First citation in article | CrossRef | ADS
Jones, B. F. 1973, A&AS, 9, 313 First citation in article | ADS
———. 1981, AJ, 86, 290 First citation in article | CrossRef | ADS
Landolt, A. 1979, ApJ, 231, 468 First citation in article | CrossRef | ADS
Makovoz, D., & Marleau, F. 2005, PASP, 117, 1113 First citation in article | CrossRef | ADS
McCarthy, M., & Treanor, P. 1964, Ric. Astron., 6, 535 First citation in article | ADS
Mermilliod, J.-C., Bratschi, P., & Mayor, M. 1997, A&A, 320, 74 First citation in article | ADS
Messina, S. 2001, A&A, 371, 1024 First citation in article | CrossRef | ADS
Meynet, G., Mermilliod, J.-C., & Maeder, A. 1993, A&AS, 98, 477 First citation in article | ADS
Oppenheimer, B., Basri, G., Nakajima, T., & Kulkarni, S. 1997, AJ, 113, 296 First citation in article | CrossRef | ADS
Patten, B., et al. 2006, ApJ, 651, 502 First citation in article | IOPscience | ADS
Pinfield, D., Hodgkin, S., Jameson, R., Cossburn, M., Hambly, N., & Devereux, N. 2000, MNRAS, 313, 347 First
citation in article | CrossRef | ADS
Pritchard, R. 1884, MNRAS, 44, 355 First citation in article | ADS
Prosser, C., Stauffer, J., & Kraft, R. 1991, AJ, 101, 1361 First citation in article | CrossRef | ADS
Queloz, D., Allain, S., Mermilliod, J.-C., Bouvier, J., & Mayor, M. 1998, A&A, 335, 183 First citation in article | ADS
Raboud, D., & Mermilliod, J.-C. 1998, A&A, 329, 101 First citation in article | ADS
Rieke, G., & Lebofsky, M. 1985, ApJ, 288, 618 First citation in article | CrossRef | ADS
Robinson, E. L., & Kraft, R. P. 1974, AJ, 79, 698 First citation in article | CrossRef | ADS
Rosvick, J., Mermilliod, J., & Mayor, M. 1992, A&A, 255, 130 First citation in article | ADS
Siess, L., Dufour, E., & Forestini, M. 2000, A&A, 358, 593 First citation in article | ADS
Skrutskie, M., et al. 2006, AJ, 131, 1163 First citation in article | IOPscience | ADS
Slesnick, C., Carpenter, J., Hillenbrand, L., & Mamajek, E. 2006, AJ, 132, 2665 First citation in article | IOPscience |
ADS
Soderblom, D. R., Jones, B. R., Balachandran, S., Stauffer, J. R., Duncan, D. K., Fedele, S. B., & Hudon, J. 1993, AJ,
106, 1059 First citation in article | CrossRef | ADS
Soderblom, D., Nelan, E., Benedict, G., McArthur, B., Ramirez, I., Spiesman, W., & Jones, B. 2005, AJ, 129, 1616 First
citation in article | IOPscience | ADS
Stauffer, J. 1980, AJ, 85, 1341 First citation in article | CrossRef | ADS
———. 1982, AJ, 87, 1507 First citation in article | CrossRef | ADS
———. 1984, ApJ, 280, 189 First citation in article | CrossRef | ADS
Stauffer, J. R., Caillault, J.-P., Gagne, M., Prosser, C. F., & Hartmann, L. W. 1994, ApJS, 91, 625 First citation in article |
CrossRef | ADS
Stauffer, J., Hamilton, D., Probst, R., Rieke, G., & Mateo, M. 1989, ApJ, 344, L21 First citation in article | CrossRef |
ADS
Stauffer, J. R., & Hartmann, L. W. 1987, ApJ, 318, 337 First citation in article | CrossRef | ADS
Stauffer, J. R., Hartmann, L. W., Soderblom, D. R., & Burnham, N. 1984, ApJ, 280, 202 First citation in article |
CrossRef | ADS
Stauffer, J., Klemola, A., Prosser, C., & Probst, R. 1991, AJ, 101, 980 First citation in article | CrossRef | ADS
Stauffer, J. R., Liebert, J., & Giampapa, M. 1995, AJ, 109, 298 First citation in article | CrossRef | ADS
Stauffer, J. R., et al. 1999, ApJ, 527, 219 First citation in article | IOPscience | ADS
———. 2005, AJ, 130, 1834 First citation in article | IOPscience | ADS
Steele, I., et al. 1995, MNRAS, 272, 630 First citation in article | ADS
Terlevich, E. 1987, MNRAS, 224, 193 First citation in article | ADS
Terndrup, D. M., Stauffer, J. R., Pinsonneault, M. H., Sills, A., Yuan, Y., Jones, B. F., Fischer, D., & Krishnamurthi, A.
2000, AJ, 119, 1303 First citation in article | IOPscience | ADS
Trumpler, R. J. 1921, Lick Obs. Bull., 10, 110 First citation in article | CrossRef | ADS
van Leeuwen, F., Alphenaar, P., & Brand, J. 1986, A&AS, 65, 309 First citation in article | ADS
van Leeuwen, F., Alphenaar, P., & Meys, J. J. M. 1987, A&AS, 67, 483 First citation in article | ADS
van Maanen, A. 1946, ApJ, 103, 289 First citation in article | CrossRef | ADS
Werner, M., et al. 2004, ApJS, 154, 309 First citation in article | IOPscience | ADS
White, R. E. 2003, ApJS, 148, 487 First citation in article | IOPscience | ADS
White, R. E., & Bally, J. 1993, ApJ, 409, 234 First citation in article | CrossRef | ADS
Descubren un nuevo ciclo del calendario Maya
Como si fuera un gran rompecabezas de 2 metros de altura por menos de 1 de ancho, el Tablero Este,
descubierto en el edificio I del Grupo XVI de Palenque, Chiapas, en 1993, dio la pista para otro gran hallazgo: Un Ciclo
Calendárico de 63 días. Así, luego de más de 1000 años, La Voz, el discurso de los antiguos Mayas plasmado en Estuco,
volvió a escucharse.
Después del trabajo de campo en tierras chiapanecas, Guillermo Bernal Romero, del Centro de Estudios Mayas
del Instituto Filológicas (IIFL) de la UNAM, (México), volvió a su cubículo y descifró el mensaje; la existencia de este
Ciclo que había pasado desapercibido en los estudios clásicos en torno al calendario.
Al hacer la reconstrucción, Bernal comprobó que el período estuvo asociado con el ritual de Taladrado de fuego,
(joch´K´ahk´), es decir, de generación por fricción de un fuego ritual dedicado al Dios Zarigüeya o Tlacuache.
El Ciclo 63 es una especie de eslabón perdido de un engranaje que faltaba. Se conocían otros de 7; 9, y 819 días.
El descubierto en Abril pasado es el resultado de multiplicar los primeros 2, y el tercero, de multiplicar 819 por 13.
Estos números no fueron un capricho de los Mayas, eran sagrados: Creían en la existencia de un Supra mundo o
región celeste con 13 niveles; de una terrestre, (la nuestra), con 7 estratos; y un inframundo con 9 niveles.
Respecto al 819, se ha propuesto que fue formulado para realizar cómputos de los Períodos Sinódicos, (Tiempo
que tarda un objeto en volver a aparecer en el mismo punto del cielo, con respecto al Sol, al observarlo desde la Tierra),
de Saturno, de 378 días, (63 X 6).
En 1993, Arnaldo González Cruz, Director del Proyecto Arqueológico Palenque, del Instituto Nacional de
Antropología e Historia, de la UNAM, descubrió entre los restos del Edificio I del grupo XVI, conjunto habitacional
sacerdotal, ubicado a un lado del Corazón Ceremonial de la ciudad, los fragmentos de lo que parecía un tablero.
Se encontraban dispersos, sepultados entre los escombros de la derruida construcción, donde los pedazos del
estuco, en el Período Clásico, en la época de K´inich Janahb´Pakal Il, “El Grande”, cubrieron las paredes de 2 pilastras.
Solo algunos cartuchos glíficos estaban pegados a las pilastras en su posición original.
Bernal Romero hizo un primer
estudio de estos fragmentos en 1998.
Allí descubrió un registro de 819 días.
En 2013 hubo una segunda revisión del
material, pero no fue hasta abril de
2014 que la restauradora Luz Lourdes
Herbert, de la Coordinación Nacional
de Conservación del Patrimonio
Cultural
del
INAH,
desplegó
completamente el material en camas de
arena.
Foto: A pesar del desarrollo de la
Epigrafía Maya, y del desciframiento
de los acontecimientos históricos o
míticos que relatan las inscripciones,
el calendario aún tiene aspectos
insospechados. Foto: UNAM)
Ya extendidos los cuadros de
escritura, se determinó que se trataba
de 2 tableros que estuvieron colocados
sobre jambas, pero las piezas estaban revueltas y no se sabía que cartuchos pertenecían a uno u otro “rompecabezas”. Eso
causó
Problemas, pero al observar con más detenimiento se pudo realizar la separación fina. Coincidían y tenían sentido.
Por ejemplo, con el dato del glifo del Dios Zarigüeya, en el extremo superior derecho del Tablero Este, se podía
saber cuántos cartuchos había tenido todo; cuatro columnas (dos dobles) y 14 filas, es decir, 56 espacios de escritura.
Además, el nombre de la deidad va acompañado de otros glifos, como el fuego, y antes de un verbo. A partir de una
esquina se reconstruyó todo, y aunque quedaron algunos huecos, donde ya no existen glifos, se pudo determinar que hubo
allí
El tablero Oeste se recuperó en un 30% y el Este en un 65%. La reconstrucción fue posible por la lógica del texto
del cómputo que posee las formulas bien conocidas de los ciclos calendáricos Mayas.
El segundo comprende una fecha absoluta de cuenta larga, que en nuestro calendario equivales al 28 de junio de
673; de esta los Mayas hicieron un cómputo hacia la fecha anterior, el 28 de mayo, 31 días antes, (habían transcurrido 11
días y un winal…), cuando se taladró el fuego, dedicado a la deidad Zarigüeya o el Tlacuache.
Esa ceremonia es muy significativa en el pensamiento mesoamericano; en la mitología, a ese animal se le
atribuye haber robado el fuego para dárselo al hombre.
Se conocía que los Mayas hacían esta ceremonia de manera sacrilizada, pero ahora se sabe que se realizaban
periódicamente cada 63 días. La comprobación del hecho se hizo en otro monumento, el dintel 29 de Yaxtilán, donde se
observó que un rito de taladro para el mismo Dios se ocurría en un lapso múltiplo de 63 con respecto al registro en
Palenque, es decir 13.230 días. (210 X 63).
Debido a que podía tratarse de una casualidad, se buscó otros registros, encontrándose al menos 8, como el Panel
2 de Laxtinich. El intervalo entre este y la fecha de Yaxchilán es equivalente a 345 ciclos de 63 días, es decir 21.735 días.
Esta periodicidad no podía ser casual, sino intencional. Además, es posible que este ciclo se haya utilizado para estimar el
período Sinódico de Saturno, que es de 378 días.
El ciclo 63 no fue registrado con frecuencia por los Mayas, lo que explica por qué paso desapercibido. No había
mucho elementos pero la reconstrucción de los tableros, en especial el Este, dio la pista para llegar a este período, que
explica cómo se construyeron otros factores numéricos tipo calendáricos.
Eric Thompson en 1943 descubrió que el 819 era resultado de la multiplicación de 3 cifras, 7; 9; y 13, hoy se
sabe que no es de manera serial, sino segmentada, es decir, 9 X 7; y luego 63 X 13.
Para los investigadores es fascinante, pues ahora saben que existen relaciones numéricas insospechadas que
delatan la existencia de otros ciclos. En otras palabras, la compleja maquinaria numérica que se creía resuelta aún no lo
está.
Revelan por qué la Luna no es una esfera perfecta
La Luna se sitúa a una
distancia media de la Tierra de
384.000 km y se aleja de ella
unos 3,8 centímetros por año.
Investigadores afirman
que el satélite sería ligeramente
achatado producto de las
primeras fuerzas de marea
ejercidas por la Tierra hace 4,4
millardos de años
Llena,
en
cuarto
creciente o menguante, la Luna,
por conocida que resulte para los
terrícolas, tiene sus misterios.
Un equipo de investigadores
propone en la revista Nature una
explicación a su forma, que no
es la de una esfera perfecta.
El satélite natural de la Tierra no es exactamente esférico, sino ligeramente achatado. La Luna presenta una
ligera hinchazón en su cara visible desde la Tierra, y otra en la cara oculta.
El equipo de Ian Garrick-Bethell, de la Universidad de California, explica la forma particular por los “efectos de
marea”, las fuerzas gravitacionales ejercidas por la Tierra durante la infancia de la Luna, hace 4,4 millardos de años.
El Sistema Solar se formó hace aproximadamente 4,5 millardos de años. Conforme al modelo que hoy es
corrientemente admitido, la Luna habría nacido de una colisión masiva padecida por la Tierra, que se acababa de formar.
Según los investigadores, las primeras fuerzas de marea ejercidas por la Tierra, que entonces estaba mucho más
cercana a la Luna, calentaron de forma desigual, según los lugares, la corteza de la Luna, cuando entonces ésta era un
océano de rocas en fusión. Este fenómeno dio a la Luna su forma, ligeramente alargada como un limón.
Más tarde, cuando la Luna se enfriaba, las fuerzas de las mareas deformaron el exterior de la Luna y fijaron sus
irregularidades.
La luna se sitúa a una distancia media de la Tierra de 384.000 km y se aleja de ella unos 3,8 centímetros por año.
Su circunferencia en el ecuador es de 10.920 km, es decir 3,7 veces inferior a la de la Tierra (40.000 km).
Scientists discover vast methane plumes escaping from Arctic seafloor
Scientists aboard the icebreaker Oden observe a methane mega flare.
Methane mega flare
event on the Laptev Sea slope of
the Arctic Ocean, at a depth of
about 62 meters. Image via
Daily Kos via University of
Stockholm.
An international team
of
scientists
aboard
the
icebreaker Oden – currently
north of eastern Siberia, in the
Arctic Ocean – is working
primarily to measure methane
emissions from the Arctic
seafloor. On July 22, 2014, only
a week into their voyage, the
team
reported
“elevated
methane levels, about 10 times
higher
than
background
seawater.” They say the culprit
in this release of methane, a
potent greenhouse gas, may be a
tongue of relatively warm water
from the Atlantic Ocean, the last
remnants of the Gulf Stream, mixing into the Arctic Ocean. A press release from University of Stockholm described the
discovery as:
… vast methane plumes escaping from the seafloor of the Laptev continental slope. These early glimpses of
what may be in store for a warming Arctic Ocean could help scientists project the future releases of the strong greenhouse
gas methane from the Arctic Ocean.
The scientists refer to the plumes as methane mega flares.
Expedition of the icebreaker Oden – called the SWERUS expedition – preliminary cruise plan and study areas of
Leg 1 and 2. EEZ=Exclusive Economic Zone; LR=Lomonosov Ridge; MR=Mendeleev Ridge; HC=Herald Canyon;
NSI=New Siberian Islands. Image via Daily Kos via University of Stockholm.
On July 22, 2014, chief scientist Örjan Gustafsson of the University of Stockholm wrote about the methane mega
flare event in his blog. He wrote:
So, what have we found in the first
couple of days of methane-focused
studies?
1) Our first observations of
elevated methane levels, about ten times
higher than in background seawater, were
documented already as we climbed up the
steep continental slope at stations in 500
and 250 meter depth. This was somewhat
of a surprise. While there has been much
speculation of the vulnerability of regular
marine hydrates [frozen methane formed
due to high pressure and low
temperature] along the Arctic rim, very
few actual observations of methane
releases due to collapsing Arctic upper
slope marine hydrates have been made. ¨
It has recently been documented that a
tongue of relatively warm Atlantic water,
with a core at depths of 200–600 meters may have warmed up some in recent years. As this Atlantic water, the last
remnants of the Gulf Stream, propagates eastward along the upper slope of the East Siberian margin, our SWERUS-C3
program is hypothesizing that this heating may lead to destabilization of upper portion of the slope methane hydrates.
This may be what we now for the first time are observing.
2) Using the mid-water sonar, we mapped out an area of several kilometers where bubbles were filling the water
column from depths of 200 to 500 meters. During the preceding 48 hours we have performed station work in two areas on
the shallow shelf with depths of 60-70m where we discovered over 100 new methane seep sites. SWERUS-C3
researchers have on earlier expeditions documented extensive venting of methane from the subsea system to the
atmosphere over the East Siberian Arctic Shelf. On this Oden expedition we have gathered a strong team to assess these
methane releases in greater detail than ever before to substantially improve our collective understanding of the methane
sources and the functioning of the system. This is information that is crucial if we are to be able to provide scientific
estimations of how these methane releases may develop in the future.
While not as long-lasting in the atmosphere as carbon dioxide, methane is much more effective than carbon
dioxide at trapping heat. Glaciologist Jason Box, in a recent and fascinating blog post (Is the climate dragon awakening?)
said: “Atmospheric methane release is a much bigger problem than atmospheric carbon dioxide release, since methane is
~20 times more powerful greenhouse gas”.
Methane has the potential to create a feedback loop in global warming. That is, as Earth’s climate warms,
methane that is frozen in reservoirs stored in Arctic tundra soils – or marine sediments – may be released into the
atmosphere. It does not last long in the atmosphere (on the order of years, rather than centuries as with carbon dixoide).
But its release will cause more warming, which will cause more methane to be released, replenishing that in the
atmosphere … causing more warming and more methane release and so on.
Methane release from the Arctic Ocean is not a new phenomenon; after all, the Stockholm scientists were there
to measure it. U.S. scientists have observed Arctic Ocean methane release, too. For example, NASA reported in April
2012 on a study in which scientists measured surprising levels of methane coming from cracks in Arctic sea ice and areas
of partial sea ice cover. In 2013, Shakova et al (2013) suggested that: … significant quantities of methane are escaping
the East Siberian Shelf as
a result of the degradation
of submarine permafrost
over thousands of years.
We suggest that bubbles
and storms facilitate the
flux of this methane to the
overlying
ocean
and
atmosphere, respectively.
Glaciologist
Jason Box, in his recent
blog post, pointed out that
methane release tends to
come in spikes, which he
calls “dragon’s breath.”
Jason
Box
published the chart above
in his blog. It shows a
possible methane spike.
Box said: “A reasonable
hypothesis for the outliers
[apparent
high
measurements of methane,
which Box calls 'dragon's
breath'] … would be:
extreme outlying positive anomalies represent high methane concentration plumes emanating from tundra and/or oceanic
sources. Another reasonable hypothesis would be: extreme outlying positive anomalies represent observational errors.
What NOAA states: the outliers ‘are thought to be not indicative of background conditions, and represent poorly mixed
air masses influenced by local or regional anthropogenic sources or strong local biospheric sources or sinks.’ Fair enough.
But the dragon breath hypothesis has me losing sleep.”
Methane bubbles discovered on Laptev continental slope of
Arctic Ocean by the science team aboard the icebreaker
Oden. Image via University of Stockholm.
On July 23, Ulf Hedman – who is aboard the Oden and who
is Science Coordinator for the Swedish Polar Research
Secretariat – gave a vivid description of the discovery in his
blog:
We are ‘sniffing’ methane. We see the bubbles on video
from the camera mounted on the CTD or the Multicorer. All
analysis tells the signs. We are in a [methane] mega flare.
We see it in the water column we read it above the surface
an we follow it up high into the sky with radars and lasers.
We see it mixed in the air and carried away with the winds. Methane in the air. Where does it come from? Is it from the
old moors and mosses that used to be on dry land but now has sunken into the sea. Does it come from the deep interior of
the Earth following structures in the bedrock up into the sand filled reservoirs collecting oil and gas then leaking out
upwards, as bubbles through the sea bed into the water, into the mid-water sonar, the Niskin bottles the analysis and into
our results?
Where does the methane come from? Is it organic or not? What’s the volume? How much is carried up into the air? Is
there an effect on the climate? One mega flare does not tell the truth. It’s not evidence enough.
We carry on for the next station.
And the next, and next, next…
Bottom line: A team of international scientists aboard the icebreaker Oden has documented “elevated methane levels,
about ten times higher than in background seawater” in the Arctic Ocean. They are calling it a methane mega flare event
and express hopes it will help them project future releases of the strong greenhouse gas methane from the Arctic Ocean,
and to understand the role this released methane might play in global warming.
Follow the SWERUS-C3 expedition – @SWERUSC3 – on Twitter.
Mystery crater in Yamal peninsula probably caused by methane release
Científicos captan olas de cinco metros en Océano Ártico
Jim Thomson, Universidad de Washington
El fenómeno sería consecuencia del calentamiento global, y acentuaría el derretimiento de hielo dentro de la
zona
Un nuevo y preocupante fenómeno ha sido
captado en el Polo Norte, específicamente, dentro del
Océano Ártico, donde no sólo se han registrado
importantes niveles de deshielo sino que ahora también, y
por primera vez, olas de cinco metros de altura.
El suceso fue captado por el experto de la
Universidad de Washington, Jim Thomson, quien detectó
durante septiembre de 2012 estas importantes olas
generadas por el viento, que no sólo serían una evidencia
del calentamiento global sino que podrían ser la causa,
además, de un derretimiento más acelerado en el hielo de
esta zona.
La investigación de Thomson muestra que
durante 2012, se llegaron a formar olas de hasta cinco metros de altura durante la parte más fuerte de una tormenta, que
habrían surgido de vientos habituales en la zona pero con una nueva realidad de mar abierto mucho más amplio en la
zona.
Este fenómeno también responde al retroceso del hielo ártico durante el verano, que sucede en un promedio
habitual de 150 kilómetros de la costa. Sin embargo, en 2012, esta cifra saltó a los 1.500 kilómetros, permitiendo una
mayor temporada de mar abierto y por tanto, un escenario más proclive a generar olas más altas.
Según los expertos, este fenómeno podría significar un importante cambio en las condiciones históricas de esta
zona de hielo y podría traer consecuencias dentro de ese ecosistema como también dentro de la navegación.
El científico junto a un grupo de otros expertos esperan evaluar más a detalle este fenómeno dentro de la zona,
con una serie de instrumentos que serán ubicados en Alaska durante los meses de verano.
Old pre-main-sequence stars
Disc reformation by Bondi-Hoyle accretion
1,2
3,4,5
4
2,4,6
P. Scicluna , G. Rosotti , J. E. Dale and L. Testi
1
ITAP, Universität zu Kiel, Leibnizstr. 15, 24118 Kiel, Germany e-mail: [email protected]
2
European Southern Observatory, Karl-Schwarzschild-Str. 85748 Garching b München, Germany
3
Max-Planck-Institut für extraterrestrische Physik, Giessenbachstraße, 85748 Garching, Germany
4
Excellence Cluster Universe, Boltzmannstr. 85748 Garching, Germany
5
Universitats-Sternwarte München, Scheinerstraße81679 München, Germany
6
INAF-Osservatorio Astrofisico di Arcetri, Largo E. Fermi 5, 50125 Firenze, Italy
Abstract
Young stars show evidence of accretion discs which evolve quickly and disperse with an e-folding time of ~3
Myr. This is in striking contrast with recent observations that suggest evidence of numerous >30 Myr old stars with an
accretion disc in large star-forming complexes. We consider whether these observations of apparently old accretors could
be explained by invoking Bondi-Hoyle accretion to rebuild a new disc around these stars during passage through a
clumpy molecular cloud. We combine a simple Monte Carlo model to explore the capture of mass by such systems with a
viscous evolution model to infer the levels of accretion that would be observed. We find that a significant fraction of stars
may capture enough material via the Bondi-Hoyle mechanism to rebuild a disc of mass ≳1 minimum-mass solar nebula,
and ≲10% accrete at observable levels at any given time. A significant fraction of the observed old accretors may be
explained with our proposed mechanism. Such accretion may provide a chance for a second epoch of planet formation,
and have unpredictable consequences for planetary evolution.
Key words: accretion, accretion disks / protoplanetary disks / circumstellar matter / stars: formation / stars: premain sequence
© ESO, 2014
1. Introduction
Circumstellar discs form around protostars as a result of angular momentum conservation during gravitational
collapse (e.g. Shu et al. 1987). In the early phases of star formation, disc material loses angular momentum and is
accreted onto the central star. The most direct observational signature of the presence of a protoplanetary disc is the
excess emission, on top of the expected naked stellar photosphere, at infrared and millimetre wavelengths, in the
ultraviolet and in optical/infrared emission lines. The long wavelength emission is produced by a dusty disc, heated by
internal dissipation processes or reprocessing of stellar radiation (e.g. Dullemond et al. 2007). The short wavelength
excess and the optical/infrared emission lines are thought to be produced by the disc-star interaction as matter accretes
onto the star or is ejected in a wind/jet (Hartmann 2009). Strong observational evidence shows that both the inner dusty
disc and accretion onto the central star quickly disappear during the early stages of pre-main-sequence evolution; the
fractions of stars with near infrared excess and with accretion signatures decay with an e-folding time of 2−3 Myr (Fedele
et al. 2010; Hernández et al. 2007). This disc dissipation timescale, even considering the possible revision by Bell et al.
(2013), sets a stringent constraint on the timescales for planet formation.
Recent work has challenged this paradigm. Sensitive, wide field Hα surveys of large star-forming complexes in
the Magellanic Clouds and our own Galaxy have revealed a population of pre-main-sequence stars that appear to be older
than 10 Myr but still show prominent Hα emission and/or infrared excess (Beccari et al. 2010; De Marchi et al. 2013a,b,
2011a,c). Although some of these “old” accretor candidates in nearby star-forming regions have been shown to be
misclassified young stellar objects (Manara et al. 2013), it is difficult to believe that this is the case for all the candidates;
these populations of old accretors are not as centrally condensed as the young stellar clusters in the same fields (e.g. De
Marchi et al. 2011b). If the line emission is interpreted as due to accretion as in young pre-main-sequence stars, the
implied accretion rates are similar to those derived at early ages, and typically higher than nearby transitional discs 1.
These findings are hard to understand in a framework in which the primordial disc is still the reservoir of accreting
material at such old ages; even one disc of age >30 Myr implies an initial population >10 5 (assuming exponential decay
with an e-folding timescale of 3 Myr).
In this paper we explore the possibility that the old accretors do not have a primordial disc, but a disc that they
re-accreted after the primordial disc had dissipated. Previous studies (Moeckel & Throop 2009; Padoan et al. 2005;
Throop & Bally 2008) have investigated the influence of Bondi-Hoyle accretion on pre-main-sequence mass-accretion
rates and the protoplanetary disc at earlier phases, during the initial evolution of the disc-star system within the progenitor
cloud. Here we investigate the possibility that a star older than 5−10 Myr happens to travel through a clumpy molecular
cloud, typically unrelated to that in which the star formed, and is able to accrete enough material to form a new accretion
disc.
2. Modelling
2.1. Bondi-Hoyle accretion
Hoyle & Lyttleton (1939), Bondi & Hoyle (1944), and Bondi (1952) proposed a mechanism by which objects
can capture matter from the interstellar medium (ISM). A massive object moving through the ISM causes a perturbation,
pulling material toward the object. As the capture of material is roughly symmetrical with respect to the direction of
motion of the star, much of the angular momentum of the material cancels out, and hence it is captured by the star to
eventually be accreted (Davies & Pringle 1980).
The rate at which material is captured is given by
(1) where v is the relative velocity between
the star and the ISM, n is the number density of the ISM, and μ is the mean molecular weight (usually taken as 2.3mH).
The gravitational cross-section is given by
, where RBH is the Bondi-Hoyle radius
the sound-speed of the ISM, typically 0.3kms-1. For a 1 M⊙ star moving at 1 km s-1, RBH ~ 1500 au.
Parameters
Values
Parameters
Values
Fv
10-2, 10-3, 10-4, 10-5
Cs
0.3 km s-1
Nstars
σ
105, 105, 106, 107
1 km s-1
Rd
α
0.1 pc
2.35
4
-3
nd
10 cm
(2)cs is
Parameters for Monte Carlo models.
To explore the effect of this process in reconstituting discs around young stars, we build a simple Monte Carlo
model to treat interactions between stars and clumps with densities typical for molecular clouds. We assume a stationary
clumpy molecular cloud, which we model as a collection of identical spherical clumps with radius Rcl and density ncl. We
parametrise the density of clumps through a volume filling factor of dense gas fV. We assume a population of “old” young
stars that has lost their primordial disc enters the cloud and moves through the clumpy medium. By randomly generating
stars with masses between 0.7 M⊙ and 3.2 M⊙2 from a Salpeter IMF (M ∝ M− αSalpeter 1955) and velocities generated
assuming a velocity dispersion of σv = 1kms-1, we sample the parameters required in Eq. (2) from the values given in
Table 1. The model simulates 10 Myr treated as a series of quasi-static time steps of length tst = 2Rcl/v∗, assuming that
each star is independent. For each star, we calculate RBH, the volume swept out per time-step
,
and hence the probability of encountering a dense clump
(3)In each time-step a uniform random
number ζ is drawn, and the star encounters a clump when ζ ≤ p; the impact parameter b of the encounter is given by
drawing a second random number ζ2 from the same generator such that
. We then determine the
accretion rate (Eq. (1)) and resolve the stellar accretion and the clump-mass depletion on a finer time-grid of 1000 substeps to accurately determine the accreted mass. Interactions where RBH>Rcl and grazing encounters are treated correctly
by taking the projected area of intersection. By repeating this process for >10 5 stars we build up meaningful statistics
about the range of possible BH accretion histories and their probabilities. Note that each star is modelled independently,
and mass accreted by a star does not influence the mass-budget available to later stars.
The accretion histories determined by this model are then passed to a viscous evolution model (Sect. 2.2) to
estimate the rate at which material is accreted by the star.
Our choice of fV is based on a reanalysis of SPH simulations of star-forming regions including feedback
mechanisms presented in Dale et al. (2012, 2013) to determine the filling factor of gas at densities higher than 10 4cm-3.
We find that for bound clouds of similar stellar mass to the regions observed by Beccari et al. (2010); De Marchi et al.
(2013b), 10-6<fV ≲ 10-3 irrespective of whether feedback from massive stars is included.
While this provides a useful estimate of the amount of mass captured in this way, it somewhat overestimates the
total as we neglect a number of physical processes. First, we neglect the motion of the clumps and assume that v = v∗ in
Eq. (2). Correct treatment of the relative motions would in general reduce RBH and hence the accretion rates. Second, stars
above 2 M⊙ have significant wind and radiation pressure that will depress the accretion rate (Edgar & Clarke 2004).
Similarly, we do not include the possible influence of the X-ray photoevaporation on the accretion, which may have an
analagous effect for lower mass stars. We also ignore the possible influence of magnetic fields, which recent studies (e.g.
Lee et al. 2014) have shown may reduce accretion rates by a factor of a few. Likewise, we neglect structure on scales
smaller than a single clump; such structure is required for a disc to form, and would reduce accretion rates relative to the
homogeneous clump case treated here. Finally, we do not include binaries. However, the only influence of binarity in the
context of Bondi-Hoyle accretion is to increase RBH, since binaries behave as a single object of mass M = M1 + M2.
Fig. 1
Cumulative fraction of the stellar
population that has accreted mass
as a function of total accreted
mass. The solid blue line indicates
a filling-factor of 10-2, the dotted
magenta line 10-3, the dashed red
line 10-4, and the dot-dashed
green line 10-5.
2.2. Viscous evolution modelling
Due to the angular momentum of the material accreted from the clump, which may be due to a density gradient
within the clump or the rotation of the clump itself, accretion cannot proceed directly onto the star (Ruffert 1997).
Therefore, the formation of a thin accretion disc is expected as the result of the viscous spreading of a thin ring. Throop &
Bally (2008) described the “buffer” effect of an accretion disc, but did not directly model it. We assume that the material
accreted from the medium circularises at a radius r0 = 0.1RBH. After a single impulse of accretion onto the disc, the
surface density is described by Σ(r) = M0/ (2π)δ(r − r0), where M0 is the deposited mass. Under the influence of an
effective viscosity ν that redistributes the angular momentum in the disc, the spreading ring solution (Lynden-Bell &
Pringle
1974)
describes
the
evolution
in
time
of
this
initial
surface
density,
(4)where ν is the kinematic
viscosity of the gas, Ω the Keplerian angular speed, I1/2 the modified Bessel function of order 1/2, λ = 2r3/2/
(3(GM∗)3/2νtr0), and we have specialized the expression for the ν ∝ r case. From this analytical solution, it is possible to
compute the mass accretion rate onto the star Ṁkernel. To derive the mass accretion rate history onto the star, we convolve
this function with the mass accretion rate history onto the disc:
(5)Given a stellar mass, the loading radius, and a law for viscosity, the evolution in time is now completely determined.
We fix the viscosity by using the well-known Shakura & Sunyaev (1973) prescription, ν = α(h/r)2r2Ω, where α is the
Shakura-Sunyaev parameter and h/r the aspect ratio of the disc. We choose typical values of α = 0.01 and h/r = 0.05(r/ 1
AU)1/4 (Armitage 2011). Operationally, we sample Eq. (4) numerically on a space and time grid. We integrate over space
to get the mass of the disc and we numerically differentiate the result to get the mass accretion rate kernel, which can be
convolved with the Bondi-Hoyle history (Sect. 2.1).
Fig. 2
Fraction of the population
that would be detected as
an old accretor at a given
time, plotted as a function
of
the
instantaneous
accretion
rate.
The
models are indicated
using the same colours
and line-styles as Fig. 1.
3. Results
Our model indicates that a fraction of the population ~ 40−50 × fV encounter dense regions and accrete more than
0.001M⊙ material by the end of the simulation (Fig. 1). The median accreted mass is typically ~0.01 M⊙, similar to the
mass of discs around young pre-main-sequence stars, with strong dependence on the stellar mass. In extreme cases,
however, more massive stars (>2 M⊙) with low v∗ that encounter several clumps can capture ≥M⊙. Our treatment of the
disc formation and evolution is probably inadequate for these extreme cases.
Converting the Bondi-Hoyle accretion into stellar accretion rates, we find Ṁ∗ ≲ 10-6M⊙yr-1 after the formation of
the disc. Owing to the assumptions inherent in our model, this rate declines from the peak as a power law as in primordial
discs.
By calculating the time each star spends accreting above a certain threshold accretion rate, one can derive a mean
time per star as a function of the threshold and hence an estimate of the fraction of the population which one expects to
observe accreting at a given time. As shown in Fig. 2, for a threshold rate of 10-8M⊙ yr-1 we typically find that the
cumulative probability is ~20fV, i.e. the fraction of a stellar population that one expects to observe as old accretors at a
given time is an order of magnitude larger than the volume filling-factor of dense clumps.
4. Discussion
Our primary goal is to assess whether the Bondi-Hoyle mechanism can contribute significantly to observations
of old accretors in regions with ongoing star formation, under a number of simple assumptions. This involves stars from a
previous star-formation episode, after their primordial discs have dispersed, interacting with a clumpy molecular cloud.
Our model indicates that up to several percent of the population passing through a region containing dense clumps may
accrete more than 0.001 M⊙ of material. Because of the factors indicated above (Sect. 2.1), the model is likely to
overestimate the total accreted mass. However, since the Bondi-Hoyle accretion is a well-understood process, the largest
sources of uncertainty derive from the parameters assumed as input to the model, and in particular the clump geometry
and filling factor, as well as the assumption that the accreted material will form a thin disc.
Our initial choice of filling factor was based on a reanalysis of the simulations of Dale et al. (2012, 2013) for
clouds similar to those observed to host old accretors. A further estimate can be obtained from the high-resolution submm maps of the 30 Dor region from Indebetouw et al. (2013). These reveal a wealth of clumpy structures, similar in scale
and density to the clumps in the Monte Carlo model used here. Assuming that the clumps are uniform spheres with an
average radius Rcl = 0.15pc and distributed in a cube whose depth is equal to the projected size of the observed region (10
× 10 × 10pc3) yields a filling factor of fV = 1.5 × 10-3, at the upper end of our parameter range.
The behaviour of the accretion disc depends strongly on the viscous timescale τν, as parametrised in terms of r0
and α. An order of magnitude change in τν has little effect on the observable old-accretor fractions at low thresholds, but
the fractions at high thresholds decline approximately in proportion to 1 /τν. For larger changes in viscosity, this also
affects the lowest thresholds explored in Sect. 3.
Since we do not include stars down to the peak of IMF (~0.3 M⊙) and Bondi-Hoyle accretion rates are ∝M2, we
may overestimate the total fraction of old accretors by a factor ~3 for the Salpeter IMF assumed here. However, Eq. (3) is
dominated by Rcl for low-mass stars, so one would expect a similar fraction of old accretors when Ṁ is a factor of 4
lower.
Comparisons between our model and the observations of old accretors are difficult, as there are no firm
constraints on the size of the old population (including non-accretors). Nevertheless, from Fig. 2 one can see that without
an unrealistically large filling factor (≫10-3) of dense clumps, the small, nearby star-forming regions are unlikely to
produce more than one old accretor, as their typical mass is a few hundred M⊙. As no old accretors have been identified
in these regions, this is consistent with our model. From the recent identification of a large (~3 × 10 3M⊙) diffuse
population with ages ≳10 Myr toward Orion (Bouy et al. 2014) one expects a few tens of reformed discs, although it is
unclear whether there is any overlap between this population and the Orion molecular clouds.
Observations of old accretors in large star-forming complexes typically detect up to several hundred such
sources in each observed region. Given the formation efficiency we have computed and our assumed filling factors, this
requires a total population at least of the order of 10 4 stars in the mass range of the observed old accretors, or ~3 × 10 4
stars correcting for the IMF, which must have passed through the regions in which the clumps are distributed. In the case
of NGC 3603, which is inferred to have a population ~104.2M⊙ (Rahman et al. 2013) and ~100 old accretors, this implies
either that the old population was significantly richer, or that fV is or was very high. The 30 Doradus region, on the other
hand, shows a similar total of old accretors, although the total population is likely ~100 times larger than NGC 3603.
Only a small fraction (1%) of the stars in 30Dor need to pass through regions containing dense clumps to produce the
observed numbers. In reality, fV will evolve with time, and it is possible that the difference we observe between these
regions may be due to 30Dor being more evolved, or having evolved more rapidly, than NGC 3603.
In our model, a significant fraction (up to several tens of percent) of stars capture enough material to form a
circumstellar disc of mass similar to primordial protoplanetary discs. This raises a number of interesting questions, such
as whether a second epoch of planet formation is possible, and how the interaction between inflowing material and an
existing planetary system might alter the accretion or the planetary evolution.
The answers to these queries depend strongly on how the inflowing material interacts with the existing system,
which we have not treated. Nevertheless, Bondi-Hoyle accretion presents a mechanism by which a new reservoir of
potentially planet-forming material may be built by up to a few percent of stars. This gives them a second chance to form
planets, from material that is potentially of different composition from the material that formed the star. Another
possibility is that these stars are already surrounded by a planetary system formed out of the primordial disc. If they
accrete new material, typically with an angular momentum different from that of the original planetary system, the
interaction of the new material and the existing planets may have a range of outcomes. Understanding the range of
possible outcomes will require detailed simulations of the accretion process and of the dynamical interactions with the
planetary systems which are beyond the scope of the present paper.
5. Conclusions
We have presented a model in which Bondi-Hoyle accretion by stars passing through dense clumps in the outer
regions of their natal molecular cloud leads to the re-formation of a circumstellar disc. As a result, these stars may
masquerade as pre-main-sequence objects due to ongoing accretion and the presence of infrared excess emission. A
significant part of the observed populations of old accretors in large star-forming regions may be explained by this
mechanism. As it may have wide-ranging consequences for the early evolution of planetary systems in rich stellar
environments, we believe that further investigation of this mechanism is warranted.
1
Although these are systematically lower mass objects.
2
Stars above ~3 M⊙ have strong winds which make a simple model inappropriate, while observations of old
accretors are incomplete for stars below 0.7−1 M⊙ depending on the distance to the observed region.
Acknowledgments
We wish to thank the anonymous referee for her/his careful reading of the text. The idea explored in this paper
came up during discussions at the ESO science days and star formation coffee as well as the Munich Star Formation
workshops. We thank the ESO Office for Science and all the institutes in the Munich area for providing a stimulating
environment. We thank P. Armitage, G. Beccari, G. Costigan, B. Ercolano, G. De Marchi, C. Manara, N. Moeckel, A.
Natta, P. Padoan, R. Siebenmorgen and S. Wolf for discussions and insights on the various aspects discussed in this
paper. P.S. is supported under DFG programme no. WO 857/10-1. G.R. acknowledges the support of the International
Max Planck Research School (IMPRS). This research was supported by the DFG cluster of excellence “Origin and
Structure of the Universe” (JED).
References
-
Armitage, P. J. 2011, ARA&A, 49, 195 [NASA ADS] [CrossRef] (In the text)
Beccari, G., Spezzi, L., De Marchi, G., et al. 2010, ApJ, 720, 1108 [NASA ADS] [CrossRef] (In the text)
Bell, C. P. M., Naylor, T., Mayne, N. J., Jeffries, R. D., & Littlefair, S. P. 2013, MNRAS, 434, 806 [NASA
ADS] [CrossRef] (In the text)
Bondi, H. 1952, MNRAS, 112, 195 [NASA ADS] (In the text)
Bondi, H., & Hoyle, F. 1944, MNRAS, 104, 273 [NASA ADS] (In the text)
Bouy, H., Alves, J., Bertin, E., Sarro, L. M., & Barrado, D. 2014, A&A, 564, A29 [NASA ADS] [CrossRef]
[EDP Sciences] (In the text)
Dale, J. E., Ercolano, B., & Bonnell, I. A. 2012, MNRAS, 424, 377 [NASA ADS] [CrossRef] (In the text)
Dale, J. E., Ercolano, B., & Bonnell, I. A. 2013, MNRAS, 430, 234 [NASA ADS] [CrossRef] (In the text)
Davies, R. E., & Pringle, J. E. 1980, MNRAS, 191, 599 [NASA ADS] (In the text)
De Marchi, G., Panagia, N., Romaniello, M., et al. 2011a, ApJ, 740, 11 [NASA ADS] [CrossRef] (In the text)
De Marchi, G., Panagia, N., & Sabbi, E. 2011b, ApJ, 740, 10 [NASA ADS] [CrossRef] (In the text)
De Marchi, G., Paresce, F., Panagia, N., et al. 2011c, ApJ, 739, 27 [NASA ADS] [CrossRef] (In the text)
De Marchi, G., Beccari, G., & Panagia, N. 2013a, ApJ, 775, 68 [NASA ADS] [CrossRef] (In the text)
De Marchi, G., Panagia, N., Guarcello, M. G., & Bonito, R. 2013b, MNRAS, 435, 3058 [NASA ADS]
[CrossRef] (In the text)
Dullemond, C. P., Hollenbach, D., Kamp, I., & D’Alessio, P. 2007, in Protostars and Planets V (Tucson:
University of Arizona Press), 555 (In the text)
Edgar, R., & Clarke, C. 2004, MNRAS, 349, 678 [NASA ADS] [CrossRef] (In the text)
Fedele, D., van den Ancker, M. E., Henning, T., Jayawardhana, R., & Oliveira, J. M. 2010, A&A, 510, A72
[NASA ADS] [CrossRef] [EDP Sciences] (In the text)
Hartmann, L. 2009, Accretion Processes in Star Formation, Second Edition (Cambridge University Press) (In the
text)
Hernández, J., Hartmann, L., Megeath, T., et al. 2007, ApJ, 662, 1067 [NASA ADS] [CrossRef] (In the text)
Hoyle, F., & Lyttleton, R. A. 1939, Proc. of the Cambridge Philosophical Society, 35, 405 [NASA ADS]
[CrossRef] (In the text)
Indebetouw, R., Brogan, C., Chen, C.-H. R., et al. 2013, ApJ, 774, 73 [NASA ADS] [CrossRef] (In the text)
Lee, A. T., Cunningham, A. J., McKee, C. F., & Klein, R. I. 2014, ApJ, 783, 50 [NASA ADS] [CrossRef] (In the
text)
Lynden-Bell, D., & Pringle, J. E. 1974, MNRAS, 168, 603 [NASA ADS] (In the text)
Manara, C. F., Beccari, G., Da Rio, N., et al. 2013, A&A, 558, A114 [NASA ADS] [CrossRef] [EDP Sciences]
(In the text)
Moeckel, N., & Throop, H. B. 2009, ApJ, 707, 268 [NASA ADS] [CrossRef] (In the text)
Padoan, P., Kritsuk, A., Norman, M. L., & Nordlund, Å. 2005, ApJ, 622, L61 [NASA ADS] [CrossRef] (In the
text)
Rahman, M., Matzner, C. D., & Moon, D.-S. 2013, ApJ, 766, 135 [NASA ADS] [CrossRef] (In the text)
Ruffert, M. 1997, A&A, 317, 793 [NASA ADS] (In the text)
Salpeter, E. E. 1955, ApJ, 121, 161 [NASA ADS] [CrossRef] (In the text)
Shakura, N. I., & Sunyaev, R. A. 1973, A&A, 24, 337 [NASA ADS] (In the text)
Shu, F. H., Adams, F. C., & Lizano, S. 1987, ARA&A, 25, 23 [NASA ADS] [CrossRef] (In the text)
Throop, H. B., & Bally, J. 2008, AJ, 135, 2380 [NASA ADS] [CrossRef] (In the text)
Temperaturas Anómalas en Julio - Agosto
NOAA
Si usted vive en el Hemisferio
Norte, las semanas pasadas han sido
extrañas. En lugares que son muy cálidos
en esta temporada, Este y Sureste de USA
y Oeste de Europa, las temperaturas han
sido tibias, mientras que zonas con
temperaturas medias en verano, Norte de
Europa y Costa del Pacífico de USA, han
sido ridículamente altas.
Records de altas temperaturas,
(35º C o más), han sido aproximadas o se
han roto en: Lituania, Polonia, Belarus,
Estonia, Latvia, y Suecia, a fines de Julio
y comienzos de agosto. Las altas
temperaturas han secado bosques y creado
incendios de vegetación en Siberia; en los
estados de Oregón, Washington, y
California, USA; y en las provincias de
British Columbia, Alberta y Territorios del
Noroeste de Canadá. Al mismo tiempo,
aire frío llegado de altas latitudes sobre
casi todo USA ha causado records de
bajas temperaturas diurnas y nocturnas
para la época en lugares tan al Sur como
Florida y Georgia. Las temperaturas
alcanzaron las de niveles de invierno en
las montañas de Tennessee.
Los mapas muestran las anomalías en temperaturas superficiales
entre Julio 27 y Agosto 03 de 2014. Se
hicieron con los datos colectados por el
satélite MODIS, y comparados con los
datos del mismo período entre 2005 y
2013.
La observación de temperaturas
por satélites alrededor del planeta se hace
por la cantidad de radiación infrarroja
emitida por la superficie terrestre,
calentada por la radiación solar en el día
y enfriada durante la noche. Estas no son
temperaturas absolutas, se refieren a la
temperatura del suelo al tocarse. Ellas
muestran cuanta temperatura esta fuera
del promedio.
Los colores rojos intensos
muestran temperaturas de hasta 10º C por
encima del promedio, los azules cuanto
más frío. Las áreas grises son zonas
donde los datos están incompletos, o
regiones cubiertas por nubes donde no se pudo medir la temperatura durante el período de observación.
Los meteorólogos ven varias posibles causas y relaciones para las ondas de calor y patrones de enfriamiento.
Sistemas de alta presión sobre Escandinavia y Norte de Rusia, así como en el Pacífico Noroeste de Norteamérica,
permitiendo que masas de aire estable construyan domos calientes que bloqueen los frentes de aire que pueden traer
cambios en vientos, precipitación, y temperaturas.
Estos patrones de bloqueo causan juntos inusuales curvaturas y serpenteos en la corriente de chorro, la cual toma
un patrón de dirección Norte – Sur, en el hemisferio Norte. La corriente de Jet mueve aire del Pacífico al Norte y calienta
el Noroeste de Canadá y Costa Oeste de USA; va al Sur desde áreas frías de Canadá al centro y Este de USA; y lleva aire
cálido del Atlántico hacia el Norte europeo. El Jet toma un Zigzag similar en el Oeste y Este de Siberia. Este patrón es
mucho más común en Invierno que en Verano.
References and Related Reading
Accuweather (2014, August 5) Heat Replaced by Storms in Central, Eastern Europe. Accessed August 7, 2014.
Accuweather (2014, July 29) Hot July for Much of Europe. Accessed August 7, 2014.
The Guardian (2014, July 17) Is global warming causing extreme weather via jet stream waves. Accessed August 7,
2014.
NASA Earth Observatory (2014, January 10) What Goes Around Comes Around.
NASA Goddard Institute for Space Studies (2012, August) The New Climate Dice: Public Perception of Climate Change.
Accessed August 7, 2014.
Weather.com (2014, August 1) July Cooldown Part Two: Polar Plunge Return. Accessed August 7, 2014.
Weather.com (2014, July 20) July Chill Brought Record Cold Temperatures. Accessed August 7, 2014.
Weather Extremes Blog, via Weather Underground (2014, August 5) First 100°F Temperature on Record in the Baltics.
Accessed August 7, 2014.
NASA Earth Observatory images by Jesse Allen, using data from the Land Processes Distributed Active Archive Center
(LPDAAC). Caption by Michael Carlowicz, with image interpretation from Bill Patzert (NASA JPL), Jason Samenow
(The Washington Post) and Linus Magnusson (European Centre for Medium-Range Weather Forecasting).
Instrument(s):
Terra - MODIS
Deflexión de la luz por el Sol.
Carlos Gil, ACA
Las teorías de la deflexión de la luz por un cuerpo masivo, provienen desde mediado del siglo XVII, cuando el
reverendo John Michell, un clérigo inglés y filósofo natural, razono de que si el sol fuera lo suficiente masivo, la luz no
podría escapar de su superficie.
El pionero de la descripción matemática de la gravedad, Sir Isaac Newton, aparentemente no escribió nada
acerca de los efectos que un cuerpo masivo tendrían sobre la luz, pero existe una nota en su tratado de óptica publicado
en 1.704, sobre las partículas de luz, las cuales podrían ser afectadas por la gravedad, de la misma manera que ocurre con
la materia ordinaria.
El primer cálculo sobre la deflexión de la luz por un cuerpo masivo fue publicado por el astrónomo alemán
Johann Georg von Soldner en 1.801. Soldner demostró que los rayos de luz de una estrella ubicada a una distancia similar
a la del sol, podrían reflectar la luz por un ángulo de acerca 0.90 segundos de arcos, o un cuarto de milésima de un grado.
Este ángulo corresponde a un diámetro aparente de un disco compacto (CD) visto desde una distancia de cerca 30
kilómetros (aproximadamente20 millas)
Los cálculos de Soldner fueron basados en las leyes del movimiento y gravitación de Newton, asumiendo que la
luz estaba constituida por partículas que movían rápidamente. Como sabemos, ni Soldener o astrónomos posteriores
intentaron verificar esta predicción, por una buena razón, realizar este experimento estaba más allá de la capacidad de los
instrumentos astronómicos en los inicios del siglo XIX.
Deflexión de la luz en teoría general de la relatividad
Un siglo después, en los inicios del siglo XX, Einstein desarrolla su teoría general de la relatividad. Einstein
calculo que la deflexión estimada por su teoría seria dos veces el valor establecido por la teoría de Newton.
Figura No. 1
La figuraNo.1, muestra la deflexión de los rayos de luz que pasan cerca de una masa esférica. Para hacer visible
este efecto, esta masa fue calculada haciéndola igual a la del sol, pero teniendo un diámetro, cincuenta mil millones de
vece menor.
De acuerdo a la teoría general de la relatividad, un rayo de luz aproximándose a un cuerpo masivo, tal como el
sol, desde su origen. Su trayectoria es deflectada tal como se observa en la figura No. 2. El valor del ángulo deflectado,
es inversamente proporcional a la distancia (Ro), del centro de masa.
Figura No. 2
La teoría general de la relatividad ofrece la siguiente ecuación general para la trayectoria de un rayo de luz,
afectado por la presencia de un cuerpo masivo tal como el sol
Las raíces de esta ecuación de segundo grado, ubican los valores del ángulo Ø1, en el segundo y tercer cuadrante
como se observa en la figura No.2 y cuyos valores son(
)y
), correspondiendo un valor total del ángulo de
deflexión de:
Cuando G, M y C toman los valores de:
, δtoma el valor de: δ =
[radianes] = 1,77 [segundos de arco]
La comprobación observacional de este valor, se realizó en la expedición efectuada en 1.919, organizada y
conducida por Eddigton, al visitar las islas Prince, ubicadas cerca del África, para presenciar y fotografiar un eclipse
total de sol, así como también tomar medidas de las estrellas alrededor del sol durante el eclipse, al respecto de este
resultado existe la siguiente anécdota:
Eddigton le comento a Einstein los resultados obtenidos sobre esta medición, y este simplemente dijo “– Lo sé,
la teoría es correcta - “. ¿Y si no se hubiese deflectado? Einstein: le respondió a Eddington, “Pues lo hubiera sentido por
el buen Dios. La teoría es correcta”.
Tres años antes de esta expedición, en
una carta dirigida a ArnoldSommerfeld,
Einstein escribía: “Usted se convencerá de la
Relatividad General una vez la haya estudiado.
Por consiguiente, no voy a decir una palabra
en su defensa”.
Figura No. 3
La figura No 3, muestra fotografía
tomada con el telescopio Kepler, de un sistema
binario, en el que se puede apreciar
perfectamente como la luz se curva a causa de
la gravedad.
Bibliografía.Mathematical Physics by Donald H.
Menzel.- Dover Publications, Inc. New York –
1.961
Introduction to Relativity by H. A.
Atwater – Pergamon Press, Oxford -1.974
A short course in General Relativity – J
Foster and J. D. Nightingale – Springer - New
York – 1.994. Nota del autor.-La fotografía
mostrada como la figura No. 3, ha sido
tomadadel artículo “La expedición de
Eddigthon,
Einstein
tenía
razón”
www.medciencia.com
Born between November 29 and December 18? Here’s
your constellation
Born somewhere between November 29 and
December 18? If so, chances are the sun passes in front of the
constellation Ophiuchus the Serpent Bearer on your birthday.
Now I can almost hear someone saying:
Wait a minute! There’s no Ophiuchus on the
horoscope page.
You are absolutely correct. That’s because Ophiuchus
is a constellation – not a sign – of the Zodiac. Follow the links
below to learn more about astrological signs versus
astronomical constellations, when and where to locate
Ophiuchus, some deep-sky treasures it contains, its mythology,
its science and more.
On a dark, moonless night, look for Ophichus above the bright ruddy star Antares. Image via Till Grednar.
Astrological signs versus astronomical constellations. The sun is in the sign Sagittarius from November 21 to
December 21. But, in the present-day sky, the sun is in front of the
astronomical constellation Ophiuchus from about November 29 to
December 18. In 2014, the sun enters the constellation Ophiuchus on
November 30 at 7:00 Universal Time (or for the U.S. Central Time
Zone: November 30, at 1:00 a.m. CST). Then the sun enters the
constellation Sagittarius on December 18, 2014, at 13:00 Universal
Time or 7:00 a.m. CST.
Whether you’re speaking about astrological signs or
astronomical constellations, the Zodiac depicts the narrow beltway of
stars on the stellar sphere through which the sun, moon and planets
travel continuously. The Zodiac runs astride the ecliptic – the sun’s
yearly pathway in front of the backdrop stars. The band of the Zodiac
extends some 8o north and south of the ecliptic, spanning a total of 16 o
in width.
The sun is said to enter the sign Sagittarius around November
21, or whenever the sun is precisely 30o west of the December solstice
point. The sun then enters the sign Capricorn on the December 21
solstice. So the sun passes through the sign Sagittarius for the month
period before and up to the December solstice, irrespective of the sun
shining in front of the constellation Ophiuchus from November 29 to
December 18.
By the way, the December solstice point moves one degree
westward in front of the zodiacal constellations – or backdrop stars –
in about 72 years. The
December solstice point will finally move
into the constellation Ophiuchus by the year 2269.
When and where to locate Ophiuchus. The best time to
observe Ophiuchus is during a Northern Hemisphere summer or a Southern Hemisphere winter. From the Northern
Hemisphere, late July and early August present this constellation high in the southern sky at nightfall and early evening.
It’s seen in the southwest sky on autumn evenings in the Northern Hemisphere.
This rather large constellation fills the area of sky to the north of the constellation Scorpius the Scorpion and to
the south of the constellation Hercules the Hero. If you’re familiar with Scorpius’ brightest star Antares, try star-hopping
to Ophiuchus from this ruddy gem of a star. The head of Ophiuchus is marked by the star Rasalhague (Alpha Ophiuchi).
Ophiuchus is joined in legend
and in the sky to the constellation of the
Serpent. If you have a dark sky, you
might find this is one constellation that
looks like what it’s supposed to be: a
big guy holding a snake. The name
Ophiuchus comes from two Greek
words meaning serpent and holding.
Can you see the Pipe Nebula a
little to the upper right of center? If not
Ophiuchus the Serpent Bearer.
Deep-sky
objects
in
Ophiuchus. On a night when the moon
is absent, take your binoculars and use
them to scan Ophiuchus, which lies near
the band of the Milky Way and so has
many deep-sky wonders. Ophiuchus
boasts of numerous globular clusters,
for example. The two easiest globular
clusters to see with ordinary binoculars are M10 and M12, as shown on the above chart. Through binoculars, they look
like faint puffs of light, but with the telescope, you begin to see these globular clusters for what they really are. They are
immense stellar cities spanning a hundred to a few hundred light-years in diameter, teeming with hundreds of thousands
of stars.
Another big deep-sky favorite is the Pipe Nebula, a vast interstellar cloud of gas and dust sweeping across about
7o of sky. At an arm’s length, that’s about the width of three to four fingers. This dark nebula resides at a distance of 600
to 700 light-years in southern Ophiuchus, and can be seen with
the unaided eye in a dark, transparent sky. The Pipe Nebula is
found due east of the star Antares, and due north of the stars
Shaula and Lesath. These two stars (but not the Pipe Nebula) are
shown on the above chart.
The Greek Asclepius or Latin Aesculapius. The
constellation Ophiuchus represents this legendary physician.
Ophiuchus in myth and star lore. In Greek sky lore, Ophiuchus
represents Asclepius – said to have been the first doctor –
always depicted holding a great serpent or snake. Depending on
how it’s used, a snake’s venom can either kill or cure. It’s said
that Asclepius concocted a potion from this snake venom, the
blood of the Gorgon monster and an unknown herb to bring the
dead back to life. This greatly alarmed the gods as it threatened
to undo the natural order of things.
As the good doctor was trying to bring Orion the
Hunter back to life, the god of the Underworld pleaded to Zeus,
the king of the gods, to reconsider the ramifications of the death
of death. Apparently his argument swayed the king of the gods.
Zeus confiscated the potion, removed Asclepius from Earth and
placed the gifted physician into the starry heavens.
We hardly know how the god of the Underworld made his
appeal. Perhaps he said only that which never lives never dies,
and that no mortal can have one without the other. The absence
of death means the absence – not the continuance – of life.
Sophocles may have expressed the myth’s inherent message
when saying:
Better to die, and sleep the never-waking sleep, than linger on and dare to live when the soul’s life is gone.
Possibly, the poet T.S. Eliot reechoed the theme of the ever-living story in his Four Quartets:
We die with the Dying
See they depart and w ego with them
We are born with the dead:
See, they return, and bring us with them.
In any event, the association with Asclepius with snakes is why we sometimes see a staff with a serpent wound around it
at doctor’s offices and hospitals, even today.
The great Johannes Kepler (1571 to 1630). The star known as Kepler’s supernova exploded in 1604, within the
boundaries of the constellation Ophiuchus.
Ophiuchus in history and science. It’s been more than 400 years since anyone has seen a supernova explosion of
a star within our own Milky Way galaxy. But in the year 1604, a supernova known as Kepler’s Supernova exploded onto
the scene, attaining naked-eye visibility for 18 months. It shone in southern Ophiuchus, not all that far from the Pipe
Nebula.
Kepler’s Supernova in 1604 came upon the heels of Tycho’s Supernova that lit up Cassiopeia in 1572. These
supernovae sent shock waves into the intelligentsia of Europe, which firmly believed in the Aristotelian notion of an
immutable universe outside the orbit of the moon. Tycho Brahe took a parallax measurement of the 1572 supernova,
proving that it could not be an atmospheric phenomenon. In fact, the supernova shone well beyond the moon’s orbit.
Shortly thereafter Kepler’s Supernova in 1604 seemed to drive home the point all over again.
Moreover, Tycho Brahe measured the distance of a comet in 1577, also finding it to be farther away than the
moon. Aristotelians wanted to believe comets were gases burning in the atmosphere, but once again, Tycho threw cold
water on the idea of Aristotle’s immutable universe.
What else can we tell you about Ophiuchus? Only that it lies in the direction to Barnard’s Star, which has caused
a gleam in the eye of many an earthly dreamer. This
relatively nearby star – only about six light-years away –
was the center of a controversy about possible planets
during the decade from 1963 to about 1973. Many
astronomers accepted a claim by Peter van de Kamp that
he had detected, by using astrometry, a perturbation in
the proper motion of Barnard’s Star consistent with its
having one or more planets comparable in mass with Jupiter.
Ultimately, that claim was refuted, and to date no planet has
been found for Barnard’s Star – nor are any expected.
Barnard’s Star, located in the
direction to the constellation
Ophiuchus. Our corner of the
universe got a little lonelier
when astronomers determined
in 2012 that that Barnard’s
Star – which is only six lightyears away – has no planets of
Earth’s size or larger in its
habitable zone. Bottom line:
The sun lies within
the
boundaries
of
the
constellation Ophiuchus the
Serpent Bearer for about two
weeks of every year, and thus
Ophiuchus is an informal
member of the Zodiac.
Astrological signs versus
astronomical constellations,
how to locate Ophiuchus,
some deep-sky treasures it
contains, plus charts and more.
Kepler 62e y 62f Planetas Acuosos
"Estos planetas no se
parecen a nada en nuestro sistema
solar. Están cubiertos con océanos
infinitos", dijo Lisa Kaltenegger,
del Instituto de Astronomía Max
Planck, que estudió los planetas.
Se trata de los dos
planetas de la estrella Kepler-62,
que se encuentra a 1200 años luz
de la Tierra, en la constelación
de Lira. Dos de sus cinco planetas,
llamados Kepler-62e y Kepler-62f,
están en la zona habitable de la
estrella, es decir, están a una
distancia de su sol que les permite
mantener la temperatura necesaria
para que exista el agua en estado
Líquido lo que es imprescindible
para la aparición de la vida.
En estos planetas hay
agua y mucha. La vida podría
existir, por tanto, pero no se sabe
si
podría
existir
alguna
civilización.
"La vida en estos planetas debería sobrevivir debajo del agua, lo que hace difícil conseguir los metales,
desarrollar la metalurgia y crear la electricidad requeridos para la existencia de una civilización", señala Kaltnegger.
"Sin embargo, los mundos podrían tener una gran belleza, con un océano azul bajo un sol de color naranja. Y
quién sabe, quizá podría existir vida lo suficientemente inteligente para desarrollar tecnología hasta un nivel que nos
sorprendería", añade Kaltnegger.
Meteorito en Nicaragua? septiembre 7, 2014
Posteado por Julio Vannini en Actualidad Astronómica
Eran las 11:04 pm del Sábado 6 de Septiembre del 2014. Me encontraba procesando unas fotos para mis álbumes
en Flickr cuando de repente las redes sociales en Nicaragua literalmente estallan. Comentarios alarmados tanto en
Facebook como en Twitter de un tremendo sonido semejante a una explosión que se hizo sentir en gran parte de la ciudad
capital, Managua y con repercusiones sísmicas también.
Por espacio de dos horas di seguimiento a las redes sociales en donde todo tipo de especulación salió a flote. No
era de extrañarse ya que un estruendo así de fuerte según los reportes, puso en vilo a casi toda Managua.
Una de las hipótesis que empezó a sonar con fuerza fue que un meteorito había caído sobre Managua. Varios
amigos empezaron a consultarme sobre el tema. Al respecto debo aclarar, que a falta de evidencia solida, cualquier cosas
que se diga no puede ser considerada como algo concreto, hasta que dicha evidencia saliera a la luz. A las consultas
hechas exteriorice mi opinión inicial: un meteorito que hubiese generado semejante estallido debió notarse en el cielo.
Como no había ningún avance en el termino noticioso, decidí dormirme y esperar la postura oficial la cual se
hizo pública después del mediodía de hoy (Domingo 7 de Septiembre). En resumen: se encontró un cráter de 12 metros de
diámetro por unos 5 de profundidad en los terrenos de la Fuerza Aérea de Nicaragua.
El experto de INETER que se presento en la televisión dio a conocer que la versión oficial de los hechos fue la
caída de un meteorito, basado en:
1. El tipo de estallido.
2. Registros en sismógrafos: uno del estallido inicial y otro al momento del impacto.
3. Sus memorias de un evento “similar” ocurrido hace tiempo atrás.
4. El cráter.
Foto tomada del sitio web de periódico Hoy, de
Managua.
Como seguramente habrán notado, el
meteorito resultante de la caída en Peekskill fue
lo suficientemente grande para sobrevivir su
paso ardiente por la atmósfera y caer a tierra sin
necesidad de crear un cráter. La gran mayoría de
rocas que se encuentran de tamaños similares no
han creado cráter alguno y se encuentran
simplemente a flor de suelo. Como habrán
notado también, el paso de esa roca fue bastante,
bastante llamativa. No fue algo de un resplandor sino
un evento que fue rastreable por muchos segundos en
el cielo. En lo personal he presenciado un par de
bólidos que se han quemado sobre la tierra, siendo el
más notable el avistado el 14 de Octubre del 2013
cuando iba rumbo a Granada.
Un recuento en Twitter del avistamiento.
Un cráter de ese tamaño (12 metros de
diámetro y 5.5 metros de profundidad) debió ser
causado por un objeto con suficiente mesa y energía
cinética. Aunque es difícil precisarlo así al aire, uno
puede pensar que el meteorito resultante debe ser
bastante grande. Personalmente mantengo contacto el
Dr. Plait quien por correo me ha sugerido un cuerpo
de al menos un metro de diámetro como el causante
de ese cráter. En otras palabras: pedazo de roca
espacial que se encuentra ahí enterrada!
Pero, realmente es eso lo que ocasiono el
estruendo y el cráter mostrado a los medios?
Antes de continuar quiero dejar bien en claro que lo siguiente es mi opinión personal, basada en los modestos
conocimientos de astronomía que poseo y los estudios realizados sobre cráteres de impacto y comentarios con otros
astrónomos. Soy un astrónomo profesional titulado como tal? No, solo soy un astrónomo amateur que le gusta mucho
investigar y tratar de encontrar la verdad de las cosas por medio de evidencias. No es de mi interés especular sobre qué
fue lo que paso anoche. Lo que planteare son las razones por las cuales no creo que sea un meteorito y porque no me
convence lo anunciado públicamente. Ustedes tienen todo el derecho de llevarme la contraria en esto si así lo desean.
Ok.
1. Un meteoro
lo
suficientemente
grande para dejar un
cráter así debió ser
visto
por
mucho
tiempo en el cielo.
Sábado por la noche
donde
miles
de
capitalinos
se
encuentran afuera de
sus casas y nadie vio
nada
salvo
un
resplandor a modo de
estallido en la zona de
carretera norte. Pongo
en referencia el video
de Plait. No existe tal
cosa
como
un
“meteoro sin estela”.
Me atrevo a agregar a esto que no se mostrado video alguno del evento (hasta el momento) que pudiese haber sido
captado por cámaras de seguridad instaladas en Managua.
2. Reporte de resplandor y no de bola de fuego.
3. Según el reporte oficial se registraron dos eventos sísmicos en los sismógrafos de INETER, uno del estallido y
otro del golpe en tierra. Mi pregunta: Como saben que esos registros se deben a eso? Con que otro dato hacen la
correlación? Si no hay video ni registro visual del evento, como saben a ciencia cierta que esos eventos sísmicos
provienen de un meteorito cayendo? Es decir: pudo ser cualquier otra cosa.
4. Si encontraron el cráter. Por qué no han excavado? Un meteorito no es algo peligroso una vez en tierra. Es
algo que se puede tocar inmediatamente. A diferencia de la creencia popular, no necesariamente debe haber indicio de
fuego o cosas quemadas, eso es puro cine. Cuando un meteorito cae, este ha sido frenado tanto por la atmósfera que viene
frió. Eso de estar esperando ayuda internacional para investigar realmente no es necesario. Estamos hablando del Ejercito
Nacional. No costaba nada llegar y excavar para sacar lo que haya estado ahí. Además, tienen expertos en bombas y
personal científico que los atiende y ninguno de ellos pudo saber cómo desenterrar un meteorito? Hasta yo que soy un
amateur sé lo que debo de hacer! Ahí adentro debería de haber una roca bien grande esperando ser rescatada. Dice una
formula muy conocida en el campo de la Física: Fuerza = Masa x Aceleración (F = m.a). Si la aceleración es tal que hace
que la velocidad sea terminal, entonces la masa debe ser muy grande para liberar la Energía (Fuerza de impacto)
suficiente para ese cráter. Además, es terreno blando. ESA PIEDRA DEBE ESTAR AHI!
5. Meteoritos de hielo. De donde sacaron eso? Para saber cuáles son los tipos de meteoritos existentes, les dejo el
siguiente enlace. Wikipedia: Tipos de meteoritos. Yo personalmente poseo algunas muestra pequeñas y he tenido la
oportunidad de estudiar otros más grandes en Boston, Massachussets. Por cierto, cuando traje esas muestras, el personal
de Aduanas me pidió que consiguiera certificación de algún tipo con el MAGFOR para asegurarse que no habría peligro
de contaminación extraterrestre. (Muchos Hombres de Negro o Evolución, por lo visto)
Y bien, estas son las razones por las cuales yo personalmente creo que lo de anoche no fue un meteorito sino otro
evento. ¿Que evento fue ese? No tengo la más mínima idea y como dije, no quiero especular al respecto. Pero si fue otra
cosa, considero como ciudadano Nicaragüense que se nos debe respetar y hablar con la verdad.
Ahora, si en realidad fue un meteorito y se muestra su extracción, análisis y composición, con gusto daré por
cancelada mi postura. Realmente estaría muy contento de que un evento poco probable como ese haya sucedido en
Nicaragua (por la pequeña extensión territorial que poseemos) y sobre todo que no haya causado pérdidas de vidas.
Estaría realmente feliz que se demostrara que estoy equivocado, pero no con palabras, sino con las evidencias
reales de la extracción de ese supuesto meteorito.
Mientras tanto, mis 5 pesos en la bolsa le van a que fue otra cosa.
Y ustedes, que opinan?
Imagen del Busto de Bolívar en el pico Bolívar desde la Hechicera a 15 Kilómetros de distancia con la Técnica de
Lucky Imaging.
Antonio Ballesteros Motín (Centro de Investigaciones de Astronomía, CIDA)
Agosto 2014, e-mail : ballesteros @cida.gob.ve
La técnica de Lucky Imaging (L.I.) que se utiliza en astrofotografía, consiste básicamente en tomar muchísimas
imágenes, cientos o mejor miles con tiempos de exposición muy cortos del orden de milisegundos para congelar la
turbulencia atmosférica en algunas tomas dependiendo del seeing del sitio y sumar las mejores imágenes de la serie. Esto
se hace con programas como el AutoStakker y RegiStax. Los profesionales suman desde el uno al cinco por ciento de las
mejores de la serie, porque tienen muchas imágenes que van desde 50.000 o máximo de un millón de cada objeto,
mientras que los aficionados tienen cientos o miles, y suman entre el 10 y el 50 % de las mejores de la serie.
En internet, el término “Lucky Imaging” se usa muy alegremente como, por ejemplo, “Imagen de la nebulosa de
Orión utilizando Lucky Imaging” y cuando uno va a la página donde está la imagen y sus datos uno encuentra lo
siguiente: se sumaron tres imágenes de 42 segundos c/u. Lo que quiere decir que simplemente es una suma de tres
imágenes con un total de 126 segundos de exposición. En otra página conseguimos, El trapecio en Orión con (L.I.):
imagen LRGB L: 300 x 1 seg. R,G,B: 150 x 1 seg. de cada color, que igualmente representa solo la suma de 750
imágenes para un total de 12,5 minutos de exposición.
El poder de L.I. es la selección de las mejores imágenes, porque si uno suma todas las imágenes, las pocas
enfocadas (atmosfera congelada) con las movidas o desenfocadas, que son la mayoría, resulta en una imagen borrosa.
Hagan una prueba con el RegiStax, tomen un video de un detalle de la luna o un planeta y sumen primero todas, después
50% y el 10% de las mejores, y procesen las tres imágenes con los wavelet del RegiStax o PhotoShop (niveles, mascara
de enfoque etc.) y compare las tres imágenes. Noten que los tiempos de exposición deben ser de milisegundos, ya que
tiempos de 42 seg. o de 1 seg. por imagen esto no es L.I.
La imagen que usaremos de referencia la vi en una página de internet, es una sola imagen tomada con una
cámara digital desde el sector de la Hechicera, probablemente hecha con un telescopio de 2.000 milímetros con un barlow
2X. Me pregunté si esto se puede mejorar con la técnica de L.I. considerando varias ventajas: el objeto esta fijo, no
necesito seguimiento con motor y se hace de día, además tengo varios telescopios y cámaras digitales y puedo hacer
pruebas con diferentes equipos. Las desventajas son que el objeto esta en el horizonte donde hay mucha turbulencia, es de
color negro, está a 15 kilómetros de distancia (ver calculo) y es de unos 80 centímetros de altura. Desde la ciudad de
Mérida es solo ocasionalmente temprano en las mañanas que el pico Bolívar (4.978 metros de altura m.s.n.m.) está
despejado y hay mucha nubosidad el resto del día.
Objeto a fotografiar, es la estatua del Libertador Simón Bolívar, en el pico Bolívar, la imagen se vería desde
atrás, vista desde la Hechicera, Mérida. Derecha: Imágenes tomadas de Internet.
Imagen tomada de una página de internet que usaremos de referencia para
compararlas con las nuevas, a la izquierda detalle con ampliación de la misma
imagen.
Los valores dados para la altura sobre el
nivel del mar del pico Bolívar en internet
van desde los 4978 hasta 5,007 metros,
siendo el valor más común el de 4978, de
tal manera que el cálculo nos da
aproximadamente 15 Kilómetros de
distancia al objeto.
Imagen Nº1, es del lugar donde se tomaron las imágenes del artículo, la terraza del CIDA en el sector de la
Hechicera (ver datos de la imagen en la tabla). En una ampliación de la misma imagen, no hay resolución suficiente para
observar el objeto en el pico Bolívar.
Imagen N2º de el
pico Bolívar (ver
datos de la imagen
en la tabla). A la
derecha, en un
detalle ampliado
de
la
misma
imagen, se nota la
necesidad
de
mayor
acercamiento para
obtener
mejor
resolución
del
objeto.
Imagen
Nº3
tomada con un
telescopio,
(ver
datos de la imagen
en la tabla). A la
izquierda
en
detalle ampliado
de
la
misma
imagen se observa
una mejora de la
resolución
respecto
a
la
imagen
de
referencia.
Imagen Nº4 fue tomada con un telescopio de 16 pulgadas de diámetro. En el foco primario se colocó una cámara
de video modelo Guppy Pro 503C, a color, de 5 megapixeles, tiempo de exposición de cada cuadro es de 10
milisegundos. Noten que cada cuadro pesa 14 Mb y durante la mayor parte del tiempo del video el sistema guarda los
cuadros en la memoria resultando en un archivo de video de 3,4 Gb. El video es lento de 2.8 FPS para un total de 87
segundos y 243 cuadros, capturado utilizando un programa el FireCapture V2.3. Con RegiStax se sumaron el 30% de los
mejores cuadros y procesado con wavelet. En el momento de la toma, había nubes altas oscureciendo el sitio, pero se nota
una mejora de la resolución. Esta cámara
tiene ROI (región de interés) y se puede
variar la resolución del CMOS, es una
especie de zoom con pérdida de resolución,
pero aumenta la velocidad del video.
La imagen Nº5 es de 1000x1000 pixeles, 10 milisegundos
por cuadro igual que el anterior, pero video más rápido 15 FPS y
un campo menor resultando en un archivo de 2,52 Gb., 904
cuadros, y una duración de 58 segundos. El objeto se ve más
cercano, sin embargo el resultado no fue satisfactorio ya que la
turbulencia se nota mucho más en el video y la imagen resultante
no se percibe mejora. La imagen Nº6, algunos campos con la
cámara Guppy Pro.
La imagen Nº 7 de 800x600 pixeles, se tomo después de varios meses de intentos para lograr el mejor video.
Ese día se tomaron cinco videos y el último fue el mejor de todos, de 11,7 milisegundos por cuadro, duración 116
segundos, 3809 cuadros, 32 FPS y un archivo final de 5,2 Gb. Se utilizó RegiStax V.6 para la suma y el procesado con los
wavelet y ajustes finales con PhotoShop, en este caso, se sumaron solo los mejores 50 cuadros (1,3% del total).
En conclusión, es evidente que la técnica de L.I. funciona, pero la resolución de la imagen depende fuertemente
de las condiciones climáticas del lugar, turbulencia atmosférica, iluminación del objeto, nubosidad, hora del día,
contaminación atmosférica (humo), etc.
Le agradezco a Johnny Cova por su ayuda y paciencia que tuvo durante varios meses que montamos un sin
número de veces el equipo, para lograr el video final.
Distancia Focal Equivalente
Si tenemos una lente de 200 milímetros y le colocamos una
cámara Nikon D700 que tiene el sensor de 35 mm, lo que se llama
comúnmente Full Frame o cámara para lentes Fx y se toma una
imagen cualquiera para calcular el aumento o el campo de la imagen,
se utiliza para el cálculo la focal del lente que es de 200 mm, por que
el factor de multiplicación es de uno pero si la imagen se tomó con
una Nikon D80, que tiene un sensor más pequeño el factor de
multiplicación es 1,5 y el lente tiene una focal equivalente de 300 mm.
Si se coloca una cámara como la Guppy Pro con un sensor mucho más
pequeño, el factor de multiplicación es de 6,08 en la máxima
resolución de la cámara y el lente es de 200 x 6,08 = 1.216 mm. En la
tabla esta el factor de multiplicación para diferentes sensores. La
mayoría de los datos se consiguen en internet.
Un error muy común en los datos de las imágenes de
planetas publicadas en internet que dicen por ejemplo la distancia
focal es de 6 metros con barlow 2x o 9 metros con barlow 3x con un
telescopio de 3 metros de distancia focal pero eso es cierto solo si el
sensor de la cámara es de formato de 35 milímetros que en la mayoría
de los casos no lo es, porque la cámara que utilizaron tiene un sensor de menor tamaño y la distancia focal es mucho
mayor de lo que dicen los datos..
Si coloco la cámara Guppy Pro en un telescopio Meade de 16 pulgadas de diámetro y una focal de 4064
milímetros véase tabla más abajo, calculada por mí, como es una hoja de Excel y es interactiva puedo cambiar la distancia
focal o el tamaño del pixels, si coloco 2800 mm, que es el Celestron 11 pulgadas que tengo en Caracas, los valores
cambian automáticamente en la tabla con una resolución de 800 x 600 pixeles (Véase tabla), con Meade 16 tengo una
distancia equivalente de 80 metros !!! Pero si uso el Celestron 11 me da un distancia focal equivalente de 55 metros,
estas distancias focales parecen grandes pero los aficionados para obtener imágenes de los planetas Marte, Júpiter,
Saturno usan normalmente entre 40 metros o más. Si uno divide el área del sensor de 35 mm (864 mm2) entre el área del
sensor de la Guppy Pro (25 mm2) nos da 36 veces más pequeña el área en la máxima resolución de 2588 x 1940, si utilizo
una resolución de 800 x 600 pixeles el área se reduce a un mas a 372 veces respecto a un sensor de formato de 35 mm.,
full frame (36mm x 24mm).
Referencia: SUMA DE IMÁGENES DIGITALES, partes I y II. El Mensajero Estelar, páginas 20 a la 28, año
37, Nº 68, Octubre-Diciembre 2013.
La desaparición de los géiseres gigantes de una luna de Júpiter desconcierta a los científicos
© NASA
Los géiseres gigantes detectados en 2013 por el
Telescopio Espacial Hubble de la NASA en Europa, una de las
lunas de Júpiter, parecen haber desaparecido, algo que ha dejado
desconcertados a los científicos.
Los enormes chorros de vapor de agua se han escondido
de la vista de los observadores de Europa, el más pequeño de los
4 satélites galileanos de Júpiter. Los géiseres detectados por
Hubble en diciembre de 2013 en las imágenes del menor de los
satélites joviales proporcionaban una oportunidad para el
descubrimiento de vida extraterrestre en el Sistema Solar.
Los investigadores sospechaban que el vapor salía de las
grietas que se abrían en el hielo debido a cambios
Provocados por las fuerzas de la Marea, cuando la Luna se alejaba de Júpiter.
De momento, los científicos son incapaces de explicar la desaparición de los géiseres. Las observaciones
posteriores del Hubble llevadas a cabo en enero y febrero de este año no mostraron signos de estas columnas de vapor de
alturas de hasta 200 km.
Según investigadores, los géiseres de Europa pueden ser esporádicos como los volcanes de la Tierra, a diferencia
de los expulsiones de vapor más o menos constantes que tienen lugar en el polo sur de Encélado, una de las lunas de
Saturno que también alberga un océano bajo la superficie.
La NASA busca obtener más datos sobre las expulsiones de agua antes del inicio de la misión espacial que a
mediados de la década de 2020 llevará a cabo la sonda Europa Clipper, que realizará múltiples vuelos sobre la luna helada
de Júpiter.
Lluvia de Estrellas de Las Geminíadas
Por. Jesús H. Otero A.
Este 13 de Diciembre ocurrirá la lluvia de estrellas más intensa e interesante del año y podrá verse desde todo el
país.
Las primeras noticias que mencionan esta lluvia de meteoros datan de los años 1860´s. La primera observación
conocida fue realizada por R. P. Greg de Manchester, Inglaterra, en 1862, cuando notó el radiante en la constelación de
Géminis. Casualmente B. V. Marsh and A. C. Twining, de USA, hizo el descubrimiento al mismo tiempo. Entre el 10 y el
12 de Diciembre. Por su parte Herschell las reportó entre Diciembre 12 y 13 de 1863. A partir de aquí empezaron a
hacerse más numerosas y el radiante fue catalogado como un radiante activo.
Normalmente los radiantes están relacionados a las órbitas de los cometas, pero no existe ningún cometa con esta
órbita conocido, en cambio si un asteroide llamado 3200 Phaeton. Las lluvias de estrellas ocurren cuando nuestro planeta
pasa a través del tubo de polvo que va dejando un cometa tras sucesivos pasos, al impactar con las finas partículas de
polvo que han ido siendo arrojadas por el cometa, se produce un fenómeno luminoso que es conocido como meteoro, o
estrella fugaz. Los asteroides no arrojan material, así que lo más probable es que 3200 Phaeton sea un cometa extinto
después de numerosos pasos por el Perihelio. Su período es muy corto, apenas 1.65 años, lo que explica su rápida
extinción
El número de meteoros observados se ha ido incrementando con el paso de los años. Hacia 1900, el radiante
producía unos 20 meteoros por hora, en los años 1950´s unos 65, en los 1980´s la taza horaria era de 85 meteoros por
hora. Pero el incremento sigue. Este año, por características orbitales especiales, se espera que puedan observarse hasta
200 meteoros por hora. Será una lluvia de estrellas fabulosa, con muchos meteoros rápidos, brillantes, y de color azul y
verde.
Se espera que para el 2050 el radiante produzca unos 200 meteoros horarios, pero a partir de aquí irá declinando
poco a poco, hasta desaparecer hacia el año 2100.
Este 13 de Diciembre, si el firmamento se nos presenta despejado, tendremos condiciones ideales para observar
esta hermosa lluvia de estrellas. Habrá Luna a partir de las 11h 30m y chocaremos contra uno de los filamentos más ricos
dejados por el extinto cometa.
Para observarlo hay que mirar después de las 10 pm hacia el Este, punto cardinal hacia donde nace el Sol. Allí
observará 3 estrellas alineadas brillantes y de color azul, este es
el Cinturón de Orión. Estas estrellas están metidas en un
rectángulo de 3 estrellas brillantes y una de brillo medio.
Proyecte una línea desde la estrella más brillante, de color azul
y llamada Rigel, a la segunda estrella más brillante del
rectángulo y de color rojo, Betelgause. Siga esta línea
prolongándola en el cielo, hasta llegar a 2 estrellas con brillo
casi idéntico, ellas son Pollux y Castor, estrellas principales de
Géminis, muy cerca de ellas está el punto de donde parecen
provenir los meteoros. No importa en qué lugar del cielo
aparezca este, si viene de esa dirección pertenece a las
Geminíadas.
Como regalo, a primeras horas de la noche la Tierra
contra un filamento importante del cometa Wirtanen y es muy
posible que se observe una lluvia de Meteoros entre Pegasus y
Piscis que estarán casi sobre nuestra cabeza luego del atardecer.
Esta lluvia de meteoros se estima que produzca entre 40 y 50
meteoros por hora, la Luna no interferirá nada con la
observación.
Esta es la mejor de todas las lluvias de meteoros que
ocurren en el año, pues son meteoros brillantes y se pueden
observar toda la noche.
Miembros de SOVAFA hemos descubierto varios
nuevos radiantes en los últimos años. Ellos son: α Cannis Majoridas A y α Cannis Majoridas B; Colúmbidas-Lepúsidas;
Vélidas; 42 Tauridas; 51 Androménidas y otros tres posibles radiantes aún por confirmar.
Si usted observa esta lluvia de estrellas, cuente cuantos meteoros observa en una hora y por favor envíeme esos
datos a la dirección o teléfono dado.
Si desea aprender ¿cómo observar esta lluvia de estrellas?, puede buscar ¿Cómo observar radiantes meteóricos
en nuestra página web: www.sovafa.com, o www.sovafa.org